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Near-Infrared Study of the Star-Forming Properties of the Rosette Complex


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NEAR INFRARED STUD Y OF THE ST AR-FORMING PR OPER TIES OF THE R OSETTE COMPLEX By CARLOS G. R OM AN-Z U NIGA A DISSER T A TION PRESENTED T O THE GRADU A TE SCHOOL OF THE UNIVERSITY OF FLORID A IN P AR TIAL FULFILLMENT OF THE REQ UIREMENTS FOR THE DEGREE OF DOCT OR OF PHILOSOPHY UNIVERSITY OF FLORID A 2006

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Cop yright 2006 by Carlos G. Rom an-Z u niga

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This w ork is dedicated to the memory of: Richard J. Elston (1961-2004) and Leonel Hern andez (1972-2001)

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A CKNO WLEDGMENTS This is the second time in my life I ha v e to write an ackno wledgment section for a thesis. Just as then, I will mak e my best ef fort to to a v oid omitting important names. Ho we v er if I do, let me say -as some sort of disclaimer that it w as not on purpose. My memory is v ery selecti v e and it tends to retain too much mo vie tri via, old jingles and bad jok es, while randomly erasing names, telephone numbers and birthdays. But seriously: I w ant to ackno wledge in a general w ay the Department of Astronomy at the Univ ersity of Florida accepting me as a student and for gi ving me a T eaching Assistantship during my rst tw o and a half years in graduate school. Also, for gi ving me an of ce to study access to a computer free photocop ying services and a marv elous w ork en vironment. I w ant to ackno wledge CON A CYT -Me xico for a fello wship that sponsored a major fraction my doctoral studies at the Uni v ersity of Florida. During 4 years CON A CYT ga v e me a li ving stipend, paid my tuition in full and paid for a signicant part of my health insurance costs. My most sincere gratitude to this e xcellent program. It is also important to mention that FLAMINGOS, the instrument we used to collect most of our data w as designed and constructed by the IR instrumentation group (PI: R. Elston) at the Uni v ersity of Florida, Department of Astronomy with support from NSF grant AST97-31180 and Kitt Peak National Observ atory Ne xt, I w ant to thank rst my supervisor Dr Elizabeth Lada. Her ef forts to direct my graduate career ha v e been enormous to say the least, and I feel that it is only f air to say that without her none of this w ould ha v e been possible. When I ask ed her for a summer job at the end of my rst year at UF she told me about this `nice little project i v

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on the Rosette Molecular Cloud, and ho w she w ould be happ y to ha v e me in char ge of it. W ell, I took the summer project, the project gre w to a thesis, and I fell in lo v e with it completely Six years later I am still w orking on it, pondering about its man y consequences, and hoping to k eep w orking on it for a while. Y ou see, I already tra v eled to the Rosette in my dreams, and I got to kno w the place well. Y es, there were man y challenges along the w ay both technical and personal, b ut I cannot b ut admire Elizabeth for al w ays being so enthusiast –and for being patient with my o wn ashes of o v er enthusiasm–, for ne v er allo wing me to gi v e up no matter ho w discouraging were the problems, and mostly for being there, sometimes as a supervisor sometimes as some sort of bask etball couch, sometimes as a psychologist, b ut most of the time as a friend. On top of that, Elizabeth generously used part of her grant mone y to tak e care of my salary and tuition requirements during my last year and my rst tw o summers. Dr Richard Elston constructed the instrument FLAMINGOS it with the skills and patience of a clockmak er and then he made sure we cared about it the same w ay he did. I remember those rst tw o years of the surv e y where the goals were still confusing and foggy our pipelines too b uggy and the piles of data o v erwhelming: Richard al w ays had a simple and clear w ay to solv e an y problem. Then he got ill and had to lea v e us, b ut I am glad I had enough time to learn from him a v ery important lesson, one about k eeping up the courage intact e v en in the most dif cult of circumstances. On another corner of the FLAMINGOS project is the ef fort and amazing stamina of Dr Nick Raines, who not only kno ws e v ery single cable and bolt of the instrument, b ut also e v ery inch of the telescope f acilities and e v ery ne dinning corner of T ucson. I w ant to thank Dr Jonathan W illiams, from the Institute for Astronomy at the Uni v ersity of Ha w aii, and former professor of Astronomy at UF He took the necesary time to direct one entire chapter of this thesis, sho wing me the w orld of radioastronomy all the w ay to the big leagues at FCRA O and IRAM. And because he is an e xpert on the Rosette, there were man y things I learned straight from his papers. v

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I w ant to thank the f aculty of the Astronomy Department for their patience, their help, their classes, their advice and their restless ef fort to inte grate what is on its w ay to be one of the best Astronomy programs in the w orld. I w ant to mention Dr Stanle y Dermott, chairman of the department, who ga v e me trust and support e v en in dif cult academic moments, Dr Francisco Re yes for being such a great supervisor as coordinator of the Astronomy Laboratory and a great e xample to follo w; Dr Raf ael Guzman for being al w ays a supporti v e friend and informal advisor and Dr Ata Sarajedini for al w ays being attenti v e to my career as Graduate Coordinator T o the secretary staf f of the Astronomy Department I wish I could dedicate a chapter of my thesis. There are no w ords to describe their help, specially with all the e xtra paperw ork that my condition of foreign student implies. I w ant to specially thank Catherine Cassidy for dedicating so much time to remind me of academic deadlines, al w ays with that e xtra wit and animosity Debra and Deborah for all that help with grants, assistanships and tra v el, Glenda for helping me with my emplo yee documentation, and also T race y and Ann, who are in dif ferent departments no w b ut al w ays were beautiful and patient and sho wed me ho w to ght the b ureocratic monsters. Joanna Le vine w as my of cemate, my project colleague, and a w onderful friend. Joanna is the kind of person that you can talk to about almost an ything, and belie v e me, in 5 years of sharing an of ce with someone you get to co v er a lot of con v ersation material. No w I really hope time will k eep us being friends and collaborators, e v en if she becomes the rst one in the w orld to combine astronomy and dance into a single, indi visible discipline and become insanely f amous. Thank you to Dr August Muench, also e x-of cemate, for his brilliant lessons on computing, scientic passion and insomnia. Man y ideas in this thesis became clear after commenting them with him. Other “ideas”, for the same reason, went safely to the garbage pail before being incorrectly stated. No w we will be of cemates again, which is a good opportunity to learn to follo w v ery high standards lik e he al w ays does. vi

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I also w ant to thank Dr Charles Lada, from the Harv ard Smithsonian Center for Astrophysics for man y useful comments and encouraging w ords about this project, some of them crucial for its completion. And no w adays for gi ving me an opportunity to w ork with him. Astronomy does not get much better than this. I w ant to thank Eric McK enzie for his patience coping with the man y v ersions of the pipeline and to enjo y observing and reducing star formation data just as he enjo ys observing and reducing those long inte gration plates of galaxy clusters. And thanks for his anecdotes about w alking across the United States, reading a thousand books and the nutritional virtues of peanut b utter I w ant to thank Dr Anthon y Gonzalez (no w adays a professor at UF), Dr Matthe w Horrobin, Dr Andrea Stolte and Dr Aaron Steinhauer all former postdoctoral fello ws, for taking our graduate student mess and helping to con v ert it into a rened scientic ef fort. At the mountain, the y mastered the man y twiggles of FLAMINGOS and the KPNO telescopes. Back at the of ce, the y were always a v ailable for questions and comments, and I cannot remember one single time when the answers were not gi v en to us with a big side of smiling: Anthon y al w ays w as inquisiti v e enough to disco v er obscure b ugs in the pipelines, b ut also patient enough to w ait for us to correct them. W ith Matthe w I learned tons too, and shared with him the e xperience of a data quality assesment trip to Boston, complete with bar hopping and all. From Aaron I will al w ays admire his neat style and enthusiasm, e v erything to him w as an opportunity to learn, which is f antastic. Oh, and his Simpsons collection k ept us a w ak e man y times at the telescope! Finally Andrea w as to put it in simple w ords, the person who sa v ed my project. She w as patient enough to understand the problems we were dealing with and then e v en more patient to x them. She e v en accepted my crazy idea of making a catalog softw are application in Supermongo and then rened the ef fort!. That is something to admire. Really vii

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I w ant to thank Noah Rashkind and Chris F oltz for being the guinea pig users of LongLe gs, with all that amazing vitality and enthusiasm. The y were also great trip pals in the Boston trip, which w as a lot of fun. Man y thanks to Bruno Ferreira and Jor ge Galle go. Bruno for instance, made possible a crucial gear of my analysis, digesting succesfully a cryptic little paper I reccomend him to read. But that is nothing compared with the man y man y great moments we shared as part of the Star F ormation Crib, and man y friendly gatherings. Jor ge, the other creature lurking in the ca v e 319 al w ays k ept me aoat with his cheerful style, along with man y discussions about mo vies, spanish rock bands and the meaning of life. No w adays these guys are so important in the Gainesville community (specially Maestro Y ogi Bruno) that I w onder ho w the city will cope with after the y lea v e. Thanks also for pizza, for the Indian Air sessions, for FILMINGOS and for man y other adv entures. I w ant to thank the man y UF graduate students that I met along the years. Names lik e Lauren, Elisha, Christos, Sue, Barbara, Jim, James, Scott, Rob and TJ might sound ancient to some people no w b ut the y were the same as the rest of us not a long ago. In a second layer are K elly Da vid, Debbie, Doug, Pimol, P aty V eera, Derrick, Bill, Catherine and Manuel among those who left. Craig, Gator e xtraordinaire is a separate case. And then all of those who k eep the jo y going on: Eric, Ana, Ileana, Naibi, Cynthia (thanks for the hospitality!), Suvrath, Mar garet (the cof fee hour angel), Michelle, Da vid, Aaron, Mik e, Ashle y V alerie, P aola, Miguel, Sung, Justin, Lauren, Curtis, Scott, Audra, Alison, Justin, Leah, Julian and Andre w T o the Me xicans in Gainesville Student Association, for being my f amily all these years. W ithout their support, adapting to a whole ne w country w ould ha v e been near to impossible. I ha v e a special mention for Julio Castro for being my lunch pal, jok e sidekick, mo vie critic partner and friend all these years. I also w ant to thank Eugenio y Milena (and her parents Soa and Fernando), V elia y Luis, Die go y Erica, Jor ge, Rocio, Juan, Hussein, Alicia, Maria Jose y Leo, Horacio y Maru, Arturo y Rosa Isela, Antonio viii

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y Roxanna, Sebastian y P aula, Nicasio y Miriam. I o we an apology for not putting the names of their children –I am already se v eral pages abo v e e xpected–. T oo man y moments we shared: meetings, barbecues, carni v als, September 15th parties, birthdays, you name it, b ut the best thing are the memories. T o my f amily for AL W A YS belie ving in me and being close to me despite distance and the time. My parents, Hector and Rosario ha v e al w ays been my greatest moti v ation, and e v erything I am no w I o we to them and their ef forts and man y sacrices. T o my brother and sister Esteban and Daniela, I w ant to say that we are all part of the same lo v e, and I am proud of you e v ery moment. I lo v e you all so much, and I need to be with you again so much... I only pray for that moment to come soon. T o belo v ed Sheikha Amina T eslima and all the members of the Al-Jerrahi community for k eeping the v ery essential light of my heart al w ays lit on. And for making time and distance in visible. Alhamdulilah. And Finally the person who I decided to share the road of life with. F abiola, can you belie v e that I am writing this for the second time around? And ho w man y w ords do I need no w to e xplain what I feel, if there are innite reasons for being thankful and happ y? No, I cannot start nor nish because you are my be ggining and my end. I lo v e you with e v ery part of me. ix

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T ABLE OF CONTENTS page A CKNO WLEDGMENTS . . . . . . . . . . . . . . . . i v LIST OF T ABLES . . . . . . . . . . . . . . . . . xiii LIST OF FIGURES . . . . . . . . . . . . . . . . . xi v KEY T O ABBREVIA TIONS . . . . . . . . . . . . . . xviii KEY T O SYMBOLS . . . . . . . . . . . . . . . . . xix ABSTRA CT . . . . . . . . . . . . . . . . . . . xx CHAPTER1 INTR ODUCTION . . . . . . . . . . . . . . . . 1 1.1 A Global Picture of Star F ormation . . . . . . . . . . 1 1.2 Moti v ations for the Study of the Rosette Comple x . . . . . . 4 2 THE R OSETTE COMPLEX IN MONOCER OS . . . . . . . . 7 2.1 Historical Perspecti v e . . . . . . . . . . . . . 7 2.2 The Rosette Neb ula and the Y oung Cluster NGC 2244 . . . . . 9 2.2.1 The Rosette Neb ula . . . . . . . . . . . . 9 2.2.2 NGC 2244 . . . . . . . . . . . . . . 10 2.2.3 Spectroscopic Studies . . . . . . . . . . . 12 2.2.4 Near Infrared Studies . . . . . . . . . . . 13 2.2.5 X-ray Studies . . . . . . . . . . . . . 14 2.3 The Rosette Molecular Cloud: Structure . . . . . . . . 14 2.3.1 CO studies . . . . . . . . . . . . . . 14 2.3.2 Interaction with the Rosette neb ula . . . . . . . . 18 2.4 The Rosette Molecular Cloud: Embedded Populations . . . . . 23 2.4.1 Dominance of Cluster F ormation in the Rosette Comple x . . 23 2.4.2 The Hypothesis of Sequential Star F ormation . . . . . 25 3 A NEAR-IR SUR VEY OF THE R OSETTE COMPLEX: OBSER V A TIONS . 27 3.1 The FLAMINGOS GMC Surv e y . . . . . . . . . . 27 3.2 Data Reduction . . . . . . . . . . . . . . . 29 3.2.1 The Data Reduction Pipeline: LongLe gs . . . . . . 29 3.2.2 The Photometry and Astrometry Pipeline: PinkP ack . . . 31 x

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3.3 Completeness of Sample . . . . . . . . . . . . 32 3.4 Positional Correction of Photometry . . . . . . . . . 35 3.5 Quality and Uniformity of the Surv e y . . . . . . . . . 40 3.6 Construction of Final Catalog . . . . . . . . . . . 44 3.6.1 Intrinsic quality: 2MASS Addendum . . . . . . . 44 3.6.2 Surv e y Area Mer ging . . . . . . . . . . . 44 3.7 Intrinsic Detection Constraints . . . . . . . . . . . 45 4 NEAR-IR SUR VEY : AN AL YSIS AND RESUL TS . . . . . . . 48 4.1 Introduction . . . . . . . . . . . . . . . . 48 4.2 Analysis . . . . . . . . . . . . . . . . . 50 4.2.1 The Nearest Neighbor Method . . . . . . . . . 50 4.2.2 Detection of Embedded Populations . . . . . . . 52 4.2.3 Infrared Excess Stars . . . . . . . . . . . 53 4.2.4 Magnitude Depth Restriction for IRX stars . . . . . . 55 4.2.5 Nearest Neighbor Analysis for Infrared Excess Stars . . . 58 4.2.6 Identication of Clusters . . . . . . . . . . 60 4.2.7 Properties of Clusters . . . . . . . . . . . 63 4.3 The Fraction of Stars in Clusters . . . . . . . . . . 69 4.3.1 Distrib ution of Sources with Respect to the Rosette Neb ula . 75 4.3.2 A Case for a Distrib uted Population? . . . . . . . 79 4.4 Discussion and Future W ork . . . . . . . . . . . 82 5 OBSER V A TIONS OF CLUSTER DENSE GAS ENVELOPES . . . . 87 5.1 Introduction . . . . . . . . . . . . . . . . 87 5.2 Observ ations and Data Reduction . . . . . . . . . . 89 5.3 Analysis and Results . . . . . . . . . . . . . 103 5.3.1 Presentation of the Data . . . . . . . . . . . 103 5.3.2 Local Extinction . . . . . . . . . . . . . 104 5.3.3 Calculation of Physical P arameters . . . . . . . . 105 5.3.4 T racer Ab undances . . . . . . . . . . . . 106 5.3.5 Gas Mass . . . . . . . . . . . . . . 108 5.3.6 Clump Sizes . . . . . . . . . . . . . . 109 5.3.7 Line W idths . . . . . . . . . . . . . . 110 5.3.8 Clump Masses . . . . . . . . . . . . . 110 5.3.9 Gas Dynamics . . . . . . . . . . . . . 114 5.4 The Embedded Stellar Population . . . . . . . . . . 116 5.4.1 Star F orming Ef ciencies . . . . . . . . . . 116 5.4.2 The Gas Stars Connection . . . . . . . . . . 117 5.5 Chemical Dif ferentiation in Cluster En v elopes . . . . . . . 119 5.6 Summary and Discussion . . . . . . . . . . . . 121 xi

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6 GLOB AL ASPECTS . . . . . . . . . . . . . . . 125 6.1 A Near Infrared Extinction Map for the RMC . . . . . . . 125 6.1.1 Moti v ations . . . . . . . . . . . . . . 125 6.1.2 Dust Extinction from Near -Infrared Colors: NICE and NICER 126 6.1.3 The Extinction Map . . . . . . . . . . . . 130 6.2 Indi vidual Extinction Cores . . . . . . . . . . . . 134 6.2.1 Identication and Estimation of Properties . . . . . . 134 6.2.2 Core Sizes . . . . . . . . . . . . . . 143 6.2.3 Core Masses . . . . . . . . . . . . . . 145 6.2.4 The Embedded Cluster Mass Function . . . . . . . 148 6.2.5 Star F ormation Ef cencies . . . . . . . . . . 153 6.3 Summary and Discussion . . . . . . . . . . . . 157 7 CONCLUSIONS AND FUTURE W ORK . . . . . . . . . . 161 7.1 Distrib ution of Y oung Stellar Populations in the Rosette Comple x . . 161 7.2 The Local En vironments of Y oung Clusters . . . . . . . 162 7.3 Extinction in the Rosette Comple x and Global Results . . . . . 163 7.4 Future W ork . . . . . . . . . . . . . . . . 164 APPENDIXA NEAR-IR SUR VEY DET AIL OF OBSER V A TIONS . . . . . . . 167 B MILLIMETER SUR VEY DET AIL OF OBSER V A TIONS . . . . . 173 C MILLIMETER SUR VEY DET AIL MUL TIP ANEL MAPS . . . . . 174 REFERENCES . . . . . . . . . . . . . . . . . . 177 BIOGRAPHICAL SKETCH . . . . . . . . . . . . . . . 183 xii

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LIST OF T ABLES T able page 2–1 Distance Estimates to the Rosette (NGC 2244) . . . . . . . . 11 4–1 Y oung Clusters Rosette Comple x . . . . . . . . . . . 64 5–1 Rele v ant Properties of Rosette Clusters . . . . . . . . . . 89 5–2 Clump properties for Rosette Clusters . . . . . . . . . . 102 5–3 Molecular Line P arameters . . . . . . . . . . . . . 106 5–4 Star F ormation Ef ciencies ( 13 CO(2-1)) . . . . . . . . . 116 6–1 Extinction Cores in the Rosette Comple x . . . . . . . . . 139 6–2 Star F ormation Ef ciencies ( A V cores) . . . . . . . . . . 157 A–1 Summary of near -IR observ ations. FLAMINGOS KPNO-2.1m . . . 168 A–1 Summary of near -IR observ ations. FLAMINGOS KPNO-2.1m . . . 169 A–1 Summary of near -IR observ ations. FLAMINGOS KPNO-2.1m . . . 170 A–1 Summary of near -IR observ ations. FLAMINGOS KPNO-2.1m . . . 171 A–2 Mean Photometric Scatter by Field . . . . . . . . . . . 172 B–1 IRAM Observ ations: Area Co v erage by T racer . . . . . . . . 173 xiii

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LIST OF FIGURES Figure page 2–1 A photograph of the Rosette Neb ula . . . . . . . . . . 7 2–2 Location of the Rosette Cloud in the Monoceros Comple x . . . . . 8 2–3 A H a vs. V -I diagram for NGC 2244 . . . . . . . . . . 13 2–4 A CO map of the Rosette Molecular Cloud . . . . . . . . . 16 2–5 Molecular Clumps in the Rosette Cloud . . . . . . . . . 17 2–6 IRAS 12 m and 1400 Mhz map o v erlay . . . . . . . . . 20 2–7 A 0.5-2 k eV Chandra image of the Rosette Comple x . . . . . . 22 2–8 Location of the Phelps & Lada clusters . . . . . . . . . . 24 3–1 UF/NO A O Rosette Comple x Surv e y Map . . . . . . . . . 30 3–2 Completeness Limits by Filter . . . . . . . . . . . . 33 3–3 Completeness by type of eld . . . . . . . . . . . . 34 3–4 Photometric Quality Areas . . . . . . . . . . . . . 37 3–5 Photometric Correction: Radial (J,H) . . . . . . . . . . 38 3–6 Photometric Correction: Radial (K) . . . . . . . . . . . 39 3–7 Photometric Correction: Colors . . . . . . . . . . . . 40 3–8 Photometric Correction: Color and Magnitude Diagrams . . . . . 41 3–9 Photometric Correction: Photometric Scatter . . . . . . . . 42 3–10 Distrib ution of Photometric Uncertainties by Filer . . . . . . . 43 3–11 Detectability of an Embedded Populations . . . . . . . . . 46 4–1 Areas of the Color -Color Diagram . . . . . . . . . . . 54 4–2 Contour le v el J H vs. H K Diagram for All stars in the Surv e y . . 56 4–3 Contour le v el J H vs. H K Diagrams Di vided by Brightness . . . 57 4–4 Nearest Neighbor Distrib utions for Bright IRX Stars . . . . . . 59 xi v

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4–5 Location of IRX Stars with Brightness K < 15 : 75 mag . . . . . 61 4–6 Identication of Clusters in the Rosette Comple x . . . . . . . 62 4–7 Distrib ution of Cluster Core and T otal Radii . . . . . . . . 65 4–8 Analysis Plots for Cluster PL01 . . . . . . . . . . . . 66 4–9 Analysis Plots for Cluster PL02 . . . . . . . . . . . . 67 4–10 Analysis Plots for Cluster PL03 . . . . . . . . . . . . 68 4–11 Analysis Plots for Cluster PL04 . . . . . . . . . . . . 69 4–12 Analysis Plots for Cluster PL05 . . . . . . . . . . . . 70 4–13 Analysis Plots for Cluster PL06 . . . . . . . . . . . . 71 4–14 Analysis Plots for Cluster PL07 . . . . . . . . . . . . 72 4–15 Analysis Plots for Cluster RLE08 . . . . . . . . . . . 73 4–16 Analysis Plots for Cluster RLE09 . . . . . . . . . . . 74 4–17 Analysis Plots for Cluster RLE10 . . . . . . . . . . . 75 4–18 Analysis Plots for Cluster NGC 2237 . . . . . . . . . . 76 4–19 Analysis Plots for Cluster NGC 2244 . . . . . . . . . . 77 4–20 Distrib ution of IRX stars as a function of distance to the Rosette Neb ula . 78 4–21 Cumulati v e Counts of IRX sources in Field 09 . . . . . . . . 80 4–22 Images of Distrib uted F ormation in Field 09 of the Surv e y . . . . . 81 5–1 Molecular Emission Maps: Cluster PL01 . . . . . . . . . 91 5–2 Molecular Emission Maps: Cluster PL02 . . . . . . . . . 92 5–3 Molecular Emission Maps: Cluster PL03 . . . . . . . . . 93 5–4 Molecular Emission Maps: Cluster PL04 . . . . . . . . . 94 5–5 Molecular Emission Maps: Cluster PL05 . . . . . . . . . 95 5–6 Molecular Emission Maps: Cluster PL06 . . . . . . . . . 96 5–7 Molecular Emission Maps: Cluster PL07 . . . . . . . . . 97 5–8 Molecular Emission Maps: Cluster RLE08A . . . . . . . . 98 5–9 Extinction in 13 CO(2-1) Map Areas (1) . . . . . . . . . . 99 xv

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5–10 Extinction in 13 CO(2-1)Map Areas (2) . . . . . . . . . . 100 5–11 Ab undance Ratios by T racer . . . . . . . . . . . . 107 5–12 Distrib ution of Clump Sizes by T racer . . . . . . . . . . 109 5–13 Distrib ution of Line W idths by T racer . . . . . . . . . . 110 5–14 Distrib ution of V irial Mass by T racer . . . . . . . . . . 111 5–15 Distrib ution of L TE Mass by T racer . . . . . . . . . . . 112 5–16 Distrib ution of V irial to L TE Mass ratios . . . . . . . . . 113 5–17 Comparison to Curv es of Binding Pressure . . . . . . . . . 114 5–18 Distrib ution of V elocity Gradients by T racer . . . . . . . . 115 5–19 Distrib ution of Star F ormation Ef ciencies . . . . . . . . . 117 5–20 Correlation Between Cluster Sizes and Emission Of fsets . . . . . 118 5–21 Ov erlap of HCO + (1-0) and CS(2-1) Emission for Cluster PL03 . . . 119 6–1 Near -Infrared Extinction Map of the Rosette Comple x . . . . . . 128 6–2 13 CO emission map of the Rosette Comple x . . . . . . . . 129 6–3 Contour Extinction Map with Cluster Positions . . . . . . . 131 6–4 Correlation between 13 CO and A V . . . . . . . . . . . 132 6–5 Distrib ution of indi vidual A V v alues between 0 and 30 mag . . . . 133 6–6 Extinction maps for Indi vidual Cores (1) . . . . . . . . . 136 6–7 Extinction maps for Indi vidual Cores (2) . . . . . . . . . 137 6–8 Extinction maps for Indi vidual Cores (3) . . . . . . . . . 138 6–9 Extinction maps for Indi vidual Cores (4) . . . . . . . . . 139 6–10 Extinction Core Proles (1) . . . . . . . . . . . . . 140 6–11 Extinction Core Proles (2) . . . . . . . . . . . . . 141 6–12 Extinction Core Proles (3) . . . . . . . . . . . . . 142 6–13 Extinction Core Proles (4) . . . . . . . . . . . . . 143 6–14 Distrib ution of Core Radii . . . . . . . . . . . . . 144 6–15 Comparison of Cluster and Core Radii . . . . . . . . . . 145 xvi

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6–16 Cluster Radii and IRXF vs. Extinction . . . . . . . . . . 146 6–17 Mean Extinction vs. Core Radii . . . . . . . . . . . . 147 6–18 Distrib ution of Extinction Core Masses . . . . . . . . . . 148 6–19 Distrib ution of Core Mass Compared to Clusters . . . . . . . 149 6–20 The Embedded Cluster Mass Distrib ution Function . . . . . . 150 6–21 Rosette Clusters in the ECMDF . . . . . . . . . . . . 151 6–22 Extinction Cores SFE vs. Cluster Radii . . . . . . . . . . 154 6–23 SFE as a Function of Distance to NGC 2244 . . . . . . . . 155 6–24 Schematic Map of the Rosette Comple x . . . . . . . . . 156 xvii

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KEY T O ABBREVIA TIONS 2MASS T w o Micron All Sk y Surv e y ECMDF Embedded Cluster Mass Distrib ution Function FCRA O Fi v e Colle ge Radio Astronomy Observ atory FITS Fle xible Image T ransport System FLAMINGOS Florida Multi-object Imaging Near -IR Grism Observ ational Spectrometer FO V Field of V ie w GMCs Giant Molecular Clouds HBL Hydrogen Burning Limit IMF Initial Mass Function IRAF Image Reduction and Analysis F acility IRX Infrared Excess ISM Interstellar Medium NNM Nearest Neighbor Method NO A O National Optical Astronomy Observ atory OB As OB Associations PMS Pre-main sequence star PSF Point Spread Function tting method R OSA T Rosetta X-Ray Satellite SFE Star F ormation Ef cienc y SQIID Simultaneous Quad Infrared Imaging De vice SSF Sequential Star F ormation ZPT Photometric calibration zero point xviii

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KEY T O SYMBOLS a vir Ratio of V irial to L TE clump mass A V V isual Extinction E(B-V) V isual Color Excess H a H a emission HII Ionized Hydrogen J ; H ; K Near Infrared Bands at 1.2, 1.6 and 2.2 m M cc Mass of star forming e xtinction core M cl us Cluster Mass M emb Mass of Embedded Stars M L T E Clump L TE mass M sc Mass of starless e xtinction core M Units of Solar Mass M vir Clump virial mass N emb Number of Embedded Stars R cc Radius of Extinction Core R cor e Cluster Core Radius R eq Cluster Equi v alent Radius R V V isual Extinction to Excess Ratio xix

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Abstract of Dissertation Presented to the Graduate School of the Uni v ersity of Florida in P artial Fulllment of the Requirements for the De gree of Doctor of Philosophy NEAR INFRARED STUD Y OF THE ST AR-FORMING PR OPER TIES OF THE R OSETTE COMPLEX By Carlos G. Rom an-Z u niga May 2006 Chair: Elizabeth A. Lada Major Department: Astronomy The Rosette Comple x is one of the most important astrophysical laboratories for the study of star formation. In this re gion we can study the interaction of an e xpanding HII re gion –impulsed by the stellar winds from the lar ge OB association NGC 2244– with a lar ge remnant molecular cloud, which is kno wn to host se v en embedded clusters. As part of a lar ge observ ational program to study the nature of young stellar populations in giant molecular clouds, we made a complete near -infrared imaging surv e y of the Rosette Comple x using the detector FLAMINGOS. This surv e y is deep enough to detect stars near the bro wn dw arf limit, impro ving considerably o v er a v ailable databases. Ho we v er gi v en the location of the Rosette Comple x at a lar ge distance from the Sun and at a latitude close to the galactic disk, the contamination of the surv e y data by eld populations is high. In order to f acilitate the detection of young populations, we combined a selection of cloud members by means of their infrared e xcess emission with a technique to detect star clusters using distances to nearest neighbors. This w ay we were able to conrm the se v en clusters pre viously identied, and to disco v er four ne w clusters. xx

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F or e v ery stellar cluster we determined for the rst time their approximate e xtensions and number of members. W e found that the fraction of stars in clusters in the Rosette Comple x is close to 87%, which is similar to other clouds lik e Orion. Ho we v er the formation of clusters in the Rosette seems to be hea vily inuenced by the interaction with the e xpanding neb ula, as e videnced by the f act that the core of the molecular cloud, coincident with the shock front of the e xpanding neb ula contains 50% of the total cluster population. The clusters in the core are also more e xtended and more populated. Our study w as complemented with a high resolution millimeter w a v elength radio surv e y of the dense gas emission around the 8 most prominent clusters in the sample. W e conrmed that all of the clusters observ ed are still embedded in what appear to be v ery compact parental clump remnants, b ut in man y cases these gaseous en v elopes are possibly becoming gra vitationally unbound, due to the partial emer gence of the young cluster stars. The dense gas maps sho w features characteristic of the interaction of clusters their local en vironment, particularly signicant of fsets of tracer emission peaks, possibly due to chemical dif ferentiation ef fects. Our near -infrared observ ations also allo wed us to construct an e xtinction map for the elds observ ed. The map sho ws an good agreement with 13 CO emission radio maps, and allo wed us to identify the main molecular cores in the comple x. Using the mass of stars in the clusters and the mass of the emission cores we calculated star formation ef ciencies, which resulted to be signicantly lar ger at the central core of the cloud. Also, e xtinction appears to be in v ersely proportional to the size of the clusters, b ut directly proportional to the fraction of IRX sources, which is suggesti v e of e v oluti v e ef fects and a rapid dispersion of the gas after clusters are formed. The cluster emer gence time scales could be similar and e v en shorter than the T T auri phase of the stars. xxi

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CHAPTER 1 INTR ODUCTION 1.1 A Global Pictur e of Star F ormation Star F ormation is one of the main puzzles in present day Astrophysics. Along the years, it has been possible to construct a relati v ely detailed picture of the physics in v olv ed in the formation of indi vidual stars ( Shu et al. 1987 ), b ut the problem of ho w to e xtrapolate that picture to e xplain the formation of lar ge groups of stars is more complicated. F or e xample, a complete model of Star F ormation should formulate correctly the necesary rates and ef cencies of formation required to populate a galaxy lik e ours, b ut also those of more or less acti v e galaxies. It w ould also need to be a sort of general scheme that could e xplain the formation of stellar populations with similar characteristics (for e xample their mass distrib utions) in completely independent en vironments. It w ould also need to unify the physics rele v ant to the prime material (interstellar clouds) and the nal product (the stars). Progress has been made, b ut no w adays Star F ormation, as a global theory still has man y untied knots. Stars form in molecular clouds, composed mainly of molecular hydrogen, which are the densest (n > 10 3 cm 3 ) and coldest (T 10 K) components of the Interstellar Medium. A signicant fraction of this molecular material e xists in the form of lar ge comple x es called Giant Molecular Clouds (GMCs), with masses of 10 4 -10 6 M and typical sizes of 10-100 pc.) GMCs are usually surrounded by e xtended en v elopes of atomic Hydrogen with typical masses of 10 6 M Practically all kno wn GMCs with distances of less than 3 kpc ha v e been forming stars during the last 10 million years and we ha v e direct e vidence for this assumption: First, man y HII re gions located at the edges of molecular clouds are being e xpanded by the winds of young, massi v e stars. By young and massi v e we understand O and B spectral 1

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2 types, with f ast nuclear b urning rates that result in lifetimes much shorter than the age of the galaxy Also, these objects are usually located in groups, called OB Associations (OB As). These associations usually ha v e spatial densities belo w the threshold for Galactic tidal disruption ( Ambartsumian 1947 ). This f act pro vides further e vidence –in this case dynamical– that star formation is recent. Second, with the aid of infrared and millimeter -w a v e detectors de v eloped in the last three decades, we are able today to see through the optically thick clouds where stars form –an impossible task for optical telescopes. This w ay we ha v e been able to observ e stars and e v en proto-stars while the y are still embedded in their parental clouds. These embedded stellar populations are e v en younger than OB As, with typical ages of 1 Myr or less. Also, from observ ations of embedded clusters in nearby GMCs ( d < 2 kpc), there is observ ational e vidence that the majority of the stars in GMCs are formed in cluster s ( Lada et al. 1991b ; Carpenter 2000 ). Moreo v er ric h clusters (100 members or more) clearly dominate o v er small groups, as the y contain more than 80% of the observ ed embedded stellar population ( Porras et al. 2003 ; Lada & Lada 2003 ). Unfortunately the dominance of lar ge embedded clusters in a v ailable catalogs might be slightly biased by an incompleteness at the small cluster re gimes. Among the reasons for this are: a) systematic surv e ys of molecular clouds aiming for the detection of an embedded population are rather scarce; b) searches for embedded clusters, if an y are usually limited to those zones with signposts of formation (e.g. the presence of v ery luminous infrared sources); c) surv e ys are mostly based on monotonic w a v elength counts, with poor corrections for background contamination. This w ay a v ailable surv e ys ha v e led to the spotting of only the richest clusters. It is only in a fe w cases when there is a search -either additional or separatefor lo w density groups and distrib uted embedded

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3 populations 1 The main reason is that small groups are logically much more dif cult to detect, especially if there are mostly composed of lo w mass stars (which are f ainter), with lar ge spatial distrib utions and projected against a high background of reddened sources. Embedded lo w mass stars are clearly harder to observ e because the y are intrinsically f aint. Ev en so, spectroscopic studies re v eal that OB associations ha v e a much lar ger number of lo w mass than massi v e stars, in proportions that are coincident with the distrib ution or Initial Mass Function (IMF) of stars in the eld. Ho we v er the spatial density of the lo w mass component is rarely abo v e that of f aint eld stars and thus the e xact fraction of lo w mass stars in young populations is dif cult to calculate from stellar density counts alone ( Lada & K ylas 1999 ). F ortunately if stellar associations are young enough (3 Myr or less) then lo w mass stars ha v e circumstellar material that causes them to ha v e an e xcess of infrared emission, and this mak es them distinguishable from eld stars. These kinds of objects, kno wn as Line-Emission or T T auri stars, are thus a good tracer of the lo w mass component of young stellar populations, b ut their observ ation is subject to uncertainties related to the e v entual weakness of line emission, the quality of the photometry required to observ e the e xcess, and the e v entual contamination from eld stars (see Chapter 4 ). The process of formation of lo w mass stars, which leads to their coe xistence with massi v e stars is also poorly understood. Some e xistent models of cluster formation are able to account for the observ ed spatial distrib utions of stars in clusters, b ut f ail to match the observ ed physical conditions of dense cores where clusters form ( Bonnell et al. 1997 ), or do not t the number distrib ution of the observ ed IMF ( Mouscho vias 1 The term distrib uted refers to stars for which their formation process cannot be directly associated to a group or a cluster F or e xample, it could refer to stars formed in isolation or stars originally formed in a group b ut dispersed to the point that the group is no longer distinguishable (e.g. Li et al. 2002 ; G omez et al. 1993 ; Carpenter 2000 .)

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4 1991 ). Also, dense molecular cores are e xpected to e xperience signicant fragmentation prior to condensation of proto-stars, a process that is not completely understood either The current hypothesis is that mar ginally stable cores e xperience cooling via dissipation of magnetohydrodynamic turb ulence in highly e xtinguished cores ( Myers 1998 ), which leads to the fragmentation of the core into a matrix of molecular k ernels. The k ernels will end up forming stars of dif ferent masses via competiti v e accretion, with the most massi v e stars either forming closer to the center of the core where accretion rates are higher or from initially lar ger k ernels ( Bonnell et al. 2001 ). In this scheme lo w mass stars will form preferentially in the outer parts of cores, resulting in a primordial mass se gre gation. The puzzle of the global properties of star formation in GMCs, with a complete understanding of the mechanisms that lead to the dominion of lar ge groups and the for mation and role of lo w mass stars, can only be solv ed by studying GMCs in a systematic approach. This means that entire GMCS should be observ ed one at a time, with instruments po werful enough to detect lo w mass stars. Also we need to co v er as much area of the cloud as possible, independently of the presence of rich clusters signposts, so that lo w density populations if an y can also be tak en into account. In such surv e ys, we could obtain unbiased statistics of the embedded stellar populations, and it w ould be easier to form a global picture of the stellar birth phenomenon. 1.2 Moti v ations f or the Study of the Rosette Complex The Rosette Comple x is a giant star forming re gion where a v ery lar ge OB A, NGC 2244 has formed. This OB A is e v acuating the center of its original cloud by means of a po werful ionization front created by the winds of its members. At the southeast edge of this re gion, there is a lar ge molecular cloud, where se v eral embedded clusters ha v e been detected. These characteristics mak e the Rosette an e xcellent laboratory to in v estigate the properties of v ery young stellar populations. The re gion has been studied e xtensi v ely in terms of its main features, the Rosette Neb ula, its OB association

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5 NGC 2244 and the physical properties of the molecular cloud. Ho we v er the characteristics of the embedded populations ha v e been studied only to a v ery supercial le v el and it is unkno wn if there are other clusters, if the y share the cloud with a lo w density population and more important, what is their relation to the prominent NGC 2244. The molecular cloud and the neb ula appear to be in clear interaction, and a basic question is ho w the formation of the ne w clusters is related to this interaction. One approach to this problem, for e xample, w ould be to study the properties of clusters as a function of distance to the neb ula and see if an y signicant dif ferences arise, which w ould be proof of the inuence of the OB association in the ne w episode of formation occuring in the cloud. One of the main goals of this thesis is to determine, to the best possible le v el, the total number of young stars in the Rosette Comple x, as well as their distrib ution, and relati v e properties. The core of the thesis is a ne w near infrared surv e y of the re gion made with the instrument FLAMINGOS, de v eloped at the Uni v ersity of Florida, which can detect stars in the Rosette do wn to the lo w mass re gimes –a task that has not been acomplished yet. After separating from the catalog the best candidates for young stars, we apply a technique based on the calculation of local surf ace densities in order to determine the location and e xtent of the kno wn clusters. The selection of stars by their infrared e xcess –determined with the use of near -infrared colors, impro v es the cluster detection techniques used in single band photometry studies for other clouds. In the rst chapter of this thesis we mak e a re vie w of the pre vious studies of the Rosette Comple x re gion. The re vie w follo ws a roughly historical line, and ends with the fe w embedded population studies done pre vious to this w ork, moti v ating the necessity for our ne w observ ations. The second chapter of this thesis is dedicated to the description of our Rosette Comple x near -infrared surv e y detailing our observ ations, data reduction methods, and data quality assessments.

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6 The third chapter describes the analysis applied to the photometric catalogs resultant from the surv e y From this analysis we attempt to impro v e the discussion about the distrib ution of star formation in the RMC. The fourth chapter describes a complementary millimetric radio w a v e study of eight RMC clusters, which has the goal of discussing the interaction between embedded star clusters and the remnants of their parental cores. The fth chapter describes the use of near -infrared colors of stars to create a detailed e xtinction map of the Rosette Cloud, which allo ws us to compare some properties of the clusters with those of their forming cores. W e also include a rst approach to the calculation of the cluster masses, which allo ws us to study star forming ef ciencies in the comple x. Finally we present a summary of the results of the thesis and a discussion on future w ork.

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CHAPTER 2 THE R OSETTE COMPLEX IN MONOCER OS 2.1 Historical P erspecti v e The Rosette Comple x (l=207.0, b=-2.1) is located at the anticenter of the galactic disk in the constellation of Monoceros. The re gion is v ery popular partly because of the staggering beauty of its main feature: a v ery e xtended emission neb ula which hosts a lar ge central HII re gion, e v acuated by the winds of a central OB association (see Figure 2.1 ). Figure 2–1: The Rosette Neb ula. Credit: Canada-France-Ha w aii T elescope / 2003 The comple x is part of a much lar ger structure kno wn as the Northern Monoceros Re gion. This re gion comprises the Mon OB1 Cloud (host of NGC 2264 and the Cone 7

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8 Neb ula), the Monoceros Loop, and the Mon OB2 Cloud in which the Rosette is one of the most prominent features (see Figure 2–2 ). Figure 2–2: The Location of the Rosette Molecular Cloud in the Conte xt of the Monoceros Comple x re gion, from Perez ( 1991 ). The catalog name for the Rosette can be some what confusing because it is not unique: The neb ula itself is usually cataloged as NGC 2237 or NGC 2246, (especially by amateur observ ers) although NGC 2237 originally referred to the brightest patch at its west side and NGC 2246 originally pointed to a bright zone at the eastern side. In addition, while the central cluster is usually kno wn as NGC 2244, it has also been

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9 cataloged as NGC 2239. Ho we v er this designation historically referred to the brightest star in the re gion, 12 Monocerotis. The cluster w as rst noticed by Flamsteed in the late 17th century and later reported by W illiam Herschel –who did not notice the neb ulosity– and John Herschel, who disco v ered se v eral of its most conspicuous features and reported them in his general catalog ( Herschel 1864 NGC 2239 = GC 1420). Other parts of the neb ula (NGC 2237 and NGC 2246) were reported by Swift ( 1886 ) who cataloged the object as being “pretty bright [pB], v ery v ery lar ge [vvL] and dif fuse [dif f]. ” Afterw ards, the re gion w as formally kno wn as the “Swift Neb ula, ” until the name “Rosette” became more popular The total e xtent of the Rosette w as not determined until the rst photographic plates were obtained by Barnard ( 1894 ). T w o of the rst applications of Rosette Neb ula photographic data were made by Hubble ( 1922 ) in his study of dif fuse neb ulae associated with massi v e stars, and Mink o wski ( 1949 ), who published a photographic plate study along with a rst discussion on the e xpansion of the HII re gion by the central cluster O stars and the possible e xistence of Bok glob ules. He estimated the mass of the neb ula to be 10 4 M and suggested that it could be “surrounded and probably embedded in obscuring material, ” thus proposing the e xistence of the companion molecular cloud. 2.2 The Rosette Neb ula and the Y oung Cluster NGC 2244 2.2.1 The Rosette Neb ula A tab ulation of the dif ferent methods used to determine the age of the Rosette Neb ula w as done by Ogura & Ishida ( 1981 ). These v aried from studies of the properties of the central ca vity ( Kahn & Menon 1961 ; Lask er 1966 ) to e v olutionary models of the HII re gion based on the luminosity of the stars ( Hjellming 1968 ). Other methods in v olv e time scales of radiation pressure ( Mathe ws 1966 1967 ), estimates of the formation time for dark glob ules in the neb ula ( Herbig 1974 ), and the separation of [OIII] emission lines ( Smith 1973 ). The mean v alue of all these age estimates is approximately 3 1 10 6 yr

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10 A series of studies by Celnik 1983 ; 1985 ; 1986 discussed the global physical characteristics of the Rosette Comple x. The rst tw o of these are dedicated to the neb ula, while the third one is a model of the interaction with the molecular cloud. In the rst of the articles, he presented a map of the H a emission in the neb ula re gion, and calculated a total inte grated ux density of 5 10 11 W m 2 within 60 0 from the center of the HII ca vity He suggested that the emission is contained in a more or less symmetric ring with a peak at 16 0 from the center In the second paper Celnik reported radio continuum observ ations (1410 and 4750 MHz) from which he w as able to determine that the neb ula is bound by ionization, forming a spherical shell with radius of 40 pc (about 85 0 ) and a total ionized matter mass of 2.3 10 4 M Using the H112 a and He112 a recombination lines (4619 and 4621 Mhz) Celnik calculated a He + ab undance of 0.12 0.03 and a non-L TE electron temperature for the neb ula of T e = 5800 700 K – almost 1100 K abo v e the L TE – with no observ able gradient with respect to the radial distance from the center Ho we v er the in v estigation by Shipman & Clark ( 1994 ) re v ealed a good t to a T r a a = 0 : 4 model for the temperature gradient in the neb ula ca vity which, interestingly could not be adjusted to the observ ed IRAS emission. Instead, the y found that this temperature gradient w as better adjusted to a = 0 : 05 for r < 47 0 and a = 0 : 2 for 47 0 < r < 65 0 2.2.2 NGC 2244 The prominent OB association that is presumed responsible for the e v acuation of the central part of the neb ula has been the subject of man y interesting studies o v er the years. The distance to this young cluster (and therefore to the entire comple x) has been estimated man y times with slightly dif ferent results. T able 2–1 is a compilation of these v alues, from which the most commonly used is 1600 to 1700 pc. Some of the rst visual photometric studies on NGC 2244 were made by Johnson ( 1962 ), who estimated the mean color e xcess in the cluster to be E(B-V) = 0 : 46 for

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11 T able 2–1: Distance Estimates to the Rosette (NGC 2244) Author V alue (pc) Method Johnson (1962) 1660 Photoelectric Photometry Ogura & Ishida (1981) 1420 V isual Photometry P erez et.al (1987) 1670 V isual Photometry P ark & Sung (2002) 1660 V isual Photometry Hensber ge et.al (2002) 1390 Spectroscop y R V = A V = E ( B V ) = 3 : 0. This w as conrmed by T urner ( 1976 ) and later Ogura & Ishida ( 1981 ), who suggested a v alue of R = 3 : 2 0 : 15. Ogura & Ishida ( 1981 ) also proposed an age of 4 1 Myr and a star formation ef cenc y of 22% for the cluster Later Marschall et al. ( 1982 ) completed a proper motion study of 287 stars in the NGC 2244 area. The y conrmed membership for 113 objects, 52 of them from the list of Ogura & Ishida ( 1981 ). A study that combined photometry as well as spectroscop y w as completed by Perez et al. ( 1987 ). The y found that some members of NGC 2244 presented anomalous v alues of R, possibly suggesting the coe xistence of main sequence stars with v ery young objects –lik ely T T auri stars. This w as conrmed with a uvby b photometry study by Perez et al. ( 1989 ), in which 4 members presented e vidence of being true pre-main sequence (PMS) objects. The y also conrmed the age of NGC 2244 to be belo w 4 Myr b ut spread to w ards younger v alues, thus conrming a model of continuous formation. A study of great importance w as performed by P ark & Sung ( 2002 ). The y obtained UBVI and H a photometry for the cluster The y were able to determine membership for a total of 30 cluster sources and to e xtend the list of kno wn PMS candidates to 21. The y subsequently identied members coincident with R OSA T point sources catalogs and spectral types from V erschueren ( 1991 ) (see section 2.2.3 ). Six of the PMS candidates were conrmed as X-ray sources. In Figure 2–3 we sho w the P ark & Sung ( 2002 ) relation between alpha emission and V -I color for NGC 2244. (the relationship actually

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12 w ould hold for an y optical or infrared color). In this gure, PMS stars are clearly located abo v e the main sequence. Later using e v olutionary models P ark & Sung ( 2002 ) sho wed that most of the PMS stars and PMS candidates in their sample appear to ha v e masses close to 1 M and an approximate mean age of 0.4 to 0.9 Myr Because the y also estimated the main sequence turn-of f age of the cluster to be 1.9 Myr the y sho wed that the cluster has not stopped forming stars yet. Another important calculation in this article is the Initial Mass Function (IMF) of NGC 2244. The y found it has a at ( G =-0.7) IMF slope in the range 0 : 5 l o gm 2 : 0. By comparing directly to the IMF model of Scalo and to the observ ed mass function of NGC 2264, the y demonstrated that NGC 2244 is highly dominated by massi v e stars, thus conrming its status as a giant OB association. 2.2.3 Spectr oscopic Studies The most complete spectroscopic study of NGC 2244 w as done by V erschueren ( 1991 ), (see section 2.2.2 ) and it has been widely used in the literature. In particular P ark & Sung ( 2002 ) used data from this study to identify the spectral types of candidate T T auri stars in NGC 2244. A lo w resolution, single slit in v estigation by Hensber ge et al. ( 1998 ) of 2 members and 3 eld stars in the re gion of NGC 2244 yielded e vidence that these were chemically peculiar possibly magnetic stars. Later Hensber ge et al. ( 2000 ) performed spectroscopic analysis of the binary member V578 Mon, which resulted in an estimated distance slightly lo wer than other photometric estimates (see T able 2–1 ). The y also calculated the age of the system to be 2.3 0.2 Myr Finally Li et al. ( 2002 ) presented lo w resolution spectra for a sample of X-ray counterparts from the R OSA T PSPC surv e y (see also Gre gorio-Hetem et al. ( 1998 )). The y were able to conrm that v e sources had strong H a emission. T w o of the stars were

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13 Figure 2–3: A H a vs. V -I diagram for NGC 2244 from P ark & Sung ( 2002 ). The solid line represents a ZAMS relation while the dashed line is a selection limit. Filled triangles are PMS stars while open triangles are PMS candidates. Bright members are mark ed with dark lled circles. X symbols are X-ray sources and dots are non-members. found to be Herbig Ae/Be and tw o others had WTTS proles. These data indicate that X-rays are an ef cent tracer of young populations. 2.2.4 Near Infrar ed Studies Recent surv e ys in the near -infrared permit in v estigation of the e xtension and structure of the cluster In the study by Li ( 2005 ) the y analyzed data from the 2MASS all sk y surv e y and suggested that NGC 2244 had a second component located approximately 6.6 pc west of the core center Data from the FLAMINGOS surv e y reported in this thesis, appear to conrm the e xistence of this second association (see 4 ), which is coincident with the area originally labeled as NGC 2237. This area is particularly interesting because it contains lar ge dust structures kno wn as ”elephant trunks”, as well as other types of v ery young condensations of material which suggest v ery recent formation.

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14 2.2.5 X-ray Studies NGC 2244 is an important tar get for X-ray studies due to the interest in in v estigating the nature of massi v e stars as sources of high ener gy photons. The R OSA T Consortium observ ations yielded 34 X-ray sources in NGC 2244, with typical ener gies of 10 30 10 32 er gs s 1 Six of these X-ray sources are PMS candidates as reported by P ark & Sung ( 2002 ). Also, Ber gh ¨ ofer & Christian ( 2002 ) studied NGC 2244 R OSA T sources and found that objects with the f aintest X-ray emission ha v e v ery high X-ray to optical luminosity ratios. The y noted that the number of X-ray emitters associated with H a emission in NGC 2244 is remarkable. T ak en together these data gi v e strength to the hypothesis that man y X-ray emitters are young late type stars. 2.3 The Rosette Molecular Cloud: Structur e 2.3.1 CO studies A substantial part of the Interstellar Medium (ISM) e xists in molecular form. Molecular hydrogen (H 2 ) is stable and ab undant, b ut unfortunately is not easily detectable because H 2 has no permanent dipole moment and therefore its transition probabilities are v ery small. CO is considered instead the best tracer of molecular gas because of its high and constant ab undance in molecular hydrogen clouds. First attempts to detect CO emission associated with the Rosette neb ula were unsuccessful as the y pointed at the neb ula re gion, which is mostly composed of neutral and ionized gas. The observ ations reported by Blitz & Thaddeus ( 1980 ), which tar geted the southeast adjacent re gion of the neb ula, were the rst successful detections of molecular gas in the Rosette. Their NRA O surv e y mapped o v er 80% of the 12 CO emission in the area of the cloud with a 1 0 beam size, and yielded information about its lar ge scale distrib ution. The y estimated the angular e xtent of the cloud to be 3.5de g (98 pc at a distance of 1600 pc) and labeled the most prominent sub-structures. In the adapted map of Figure 2–4 we sho w an optical picture of the Rosette neb ula from the DSS o v erlayed with contours of

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15 12 CO from the Bell Labs maps of Blitz & Stark ( 1986 ). W e included the nomenclature of Blitz & Thaddeus ( 1980 ). P articularly important re gions are the Monocer os Ridg e (re gion A1-2) which is literally a re gion of gas compression at the cloud-neb ula interf ace; the Cloud Centr al Cor e (A1-1), which hosts the most massi v e clumps in the cloud and is the strongest re gion of star formation; the cores D and G, which are separated from the main body of the cloud b ut ha v e ongoing star formation; the IRS cor e which hosts the massi v e protobinary AFGL-961 (see §4 and §5.1); the Bac k Cor e B, which is more loose in structure than the re gions near the neb ula; and the Arm or E core, which despite its brightness contains no signicant star formation (no IRAS sources, or near -infrared clusters ha v e been found in this core so f ar). In a subsequent study Blitz & Stark ( 1986 ) mapped the 12 CO and 13 CO emission with impro v ed sensiti vity at the A T&T Bell Labs, unco v ering the high de gree of clumpiness of the cloud. The study of W illiams et al. ( 1994 ) also made use of the data from Blitz & Stark ( 1986 ) and listed a total of 95 clumps. The clumps with e vidence of star formation had lar ger peak temperatures, lar ger densities and also were more gra vitationally bound compared to clumps from the Maddalena Comple x, a cloud with v ery lo w star formation. Later W illiams et al. ( 1995 ) sho wed that about half of the clumps in the RMC were gra vitationally bound and the rest were supported by pressure from the interclump medium, which w as sho wn to be mostly atomic and about 40 times less dense. In gure 2–5 we sho w the locations and relati v e sizes of the clumps from W illiams et al. ( 1995 ). From the clump central v elocities W illiams et al. ( 1995 ) found that the cloud has a well dened v elocity gradient of about 0.08 km s 1 pc 1 Also, the ne gati v e correlation between clump mass and clump to clump v elocity dispersion suggested that the system is still f ar from equipartition e v en though it is dynamically e v olv ed. W illiams et al. ( 1995 ) also found the star forming acti vity to be more intense in the ridge and the central core areas, near the interf ace cloud-neb ula: clumps located near the neb ula presented lar ger

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16 Figure 2–4: A map of the Rosette Comple x area. The background image is a DSS plate of the IRS surv e y at 25 microns. The contours represent 12 CO inte grated intensity le v els from the surv e y of Blitz & Stark ( 1986 ). Indicated with labels are the main re gions of the molecular comple x identied by Blitz & Thaddeus ( 1980 ) e xcitation temperatures, a v erage densities and star forming ef ciencies and could be translated as rought clues of e v olution. Other properties of the clumps (mass, sizes or line widths) did not sho w an y signicant v ariations along the cloud. Another CO study w as done by Schneider et al. ( 1998a ). The observ ations focused on the central part of the cloud, detailing the structure of the midplane star forming cores. The y paid special attention to the IRS core, where the source AFGL-961 is located

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17 Figure 2–5: Locations and relati v e sizes of clumps in the Rosette Molecular Cloud from W illiams et al. ( 1995 ). The size of the symbol is proportional to the mass of the clump. (see also section §5.1) and pointed out the ample blue wing emission due to the po werful outo w from this object. In a complementary study ( Schneider et al. 1998b ), e xamined the CII emission (158 m) at the ridge, the central core and the IRS core. The y found weak b ut signicant C + emission deep into the molecular cloud cores and suggested that the distrib ution agrees well with a clump y molecular cloud e xposed to a lo w le v el UV radiation eld. The penetration of UV photons in the cloud is apparently f acilitated by a high density contrast clump-interclump medium. The clump mass spectrum in the RMC has the form d N = d M M x where x 1.6, with small v ariations in the e xponent depending on the range, bin size and beam resolution used (for e xample W illiams et al. suggest that x is closer to 1.3). The e xponent in this po wer la w is similar to other clouds ( Blitz 1993 ), b ut what is more important, it sho ws that the clump mass spectrum is much shallo wer than the observ ed stellar IMF ( x = 2 : 35). This sho ws that although small clumps ha v e a lar ger number proportion, most of the mass is contained in only a fe w big clumps, while for stars both numbers and total mass are dominated by the lo west mass bins. Interestingly enough, the po wer la w inde x

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18 is in f act similar to that corresponding to the mass distrib ution function for embedded clusters ( Lada & Lada 2003 ), which is suggesti v e of a uniform star formation ef cenc y for most star forming cores. The more recent surv e y of He yer et al. ( 2005 ), obtained with the wide eld array SEQ UOIA at the FCRA O 14m telescope ha v e a resolution of 45 00 at 115 GHz and 47 00 at 110 GHz. The maps re v eal “te xtural v ariations” in the 12 CO emission across the comple x, with a brighter emission component within the neb ula projected radius (approx. 40 pc from the center as dened by Celnik ( 1986 )) and weak er more e xtended emission outside this ionization edge. HBW05 suggest that the weak er emission is probably due to subthermally e xcited material with lo wer densities. The y also calculated the total molecular mass of the cloud to be 1 : 6 10 5 M from 12 CO, and found a L TE mass of 1 : 16 10 5 M from 13 CO. Moreo v er the y were able to apply a Principal Component Analysis ( He yer & Schloerb 1997 ) to determine the turb ulent o ws and the turb ulence scale in the RMC. This analysis re v eals more signicant v ariations in the v elocity structure of the cloud at the re gions located within the ionization than in the more dif fuse, e xternal component. This f act re v eals the interaction of the cloud and the HII re gion. The y suggested, ho we v er that these interactions are still v ery localized, and ha v e not af fected the global dynamics of the cloud yet. 2.3.2 Interaction with the Rosette neb ula In his third study of the Rosette, Celnik ( 1986 ) focused on comparing his H a map and radio continuum observ ations of the neb ula (see section §2.1) with the CO map of the molecular cloud from Blitz & Thaddeus ( 1980 ). Celnik constructed a comple x model of the distrib ution of the main CO cores (see Fig 2–4 ) in the conte xt of the HII re gion and estimated the rotation center of the cloud at ( a ; d ) = ( 98 : 1615 ; 4 : 3287 ; J 2000 : 0 ) Finally he re-calculated of the mass of the entire comple x by adding the total mass of ionized atoms, stars, dust and molecular gas, resulting in 3 : 3 10 5 M

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19 Cox et al. ( 1990 ), used the a v ailable IRAS data (12, 25, 60 and 100 m), and determined in great detail the distrib ution of dust and compared this to the distrib utions of ionized and molecular gas. Additionally the y were able to estimate a total infrared luminosity of roughly 1.1 10 6 L or about 50% of the a v ailable luminosity from the cluster NGC 2244. W arm dust (usually present near an OB association) typically emits strongly at the four IRAS bands. Ho we v er Cox et al. also sho wed that in the Rosette, while the 60 and 100 m emission were quite strong at re gions of ionized and neutral gas (neb ula), the 12 m emission w as preferently located be yond the limits of the ionization front (molecular cloud), suggesting a hea vy rate of destruction of dust grains from UV radiation from the cluster Surprisingly the 25 m emission w as found to be signicant in some parts of the ionized neb ula, possibly due to the e xistence of a second type of dust particle that is more resistant to UV photons. This w as also suggested by Shipman & Clark ( 1994 ), who found that the maximum temperature in the shallo w temperature gradient found in the neb ula (see section §2.1) seems too lo w to sublimate ice mantles in grains and too lo w for grains to emit signicantly in 12 or 25 m –a problem possibly solv ed with a second type of grain in the re gion. Later Shipman & Care y ( 1996 ) suggested that line emission could be contrib uting strongly to the IR emission of the neb ula, and suggested, once more that the presence of a ”hot dust” component is necessary to model this emission, especially for the 25 m. Figure 2–6 sho ws the superposition of the 12 m emission from the IRAS surv e y and a 1400 MHz radio continuum emission map from Holda w ay Braun & Liszt (unpublished). The infrared contours indicate that the w arm dust emission denes a shell that encloses the ionization front, sho wing the ef fect of hea vy dust destruction by the neb ula. The o v erposition of these maps denes v ery clearly the re gion where the HII re gion impacts the molecular cloud. K uchar & Bania ( 1993 ) made a complete map of HI emission at 21cm using the Arrecibo telescope. The y found that atomic gas in the Rosette Comple x is distrib uted

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20 Figure 2–6: IRAS 12 m emission map superimposed on a 1400 MHz radio continuum emission map by Holda w ay Braun and Liszt, NRA O. in three main re gions which form a rough, e xtended shell around the optical neb ula and be yond the molecular cloud, with a center of e xpansion at ( a ; d ) = (97.95,4.97, J2000). This shell (according to their calculations) w ould ha v e a mass close to 2 10 4 M which implies a b udget of kinetic ener gy for the shell e xpansion of approx. 4 10 48 er gs, or 2% of the total ener gy a v ailable from the stars in NGC 2244.

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21 High Ener gy studies The possibility of interaction between the HII re gion and the star forming cloud, as well as the location of the Rosette Neb ula near the edge of the Monoceros Loop –also kno wn as the Monoceros Superno v a Remnant– ( Da vies 1963 ), has moti v ated a number of studies aimed at in v estigating the high ener gy photon emission in the interaction re gions. Deep H a +[NII] photographic plates by Da vies et al. ( 1978 ) suggested a correlation between a lamentary structure observ able in H a emission and a Rosette neb ula feature observ able in 240 MHz radio w a v es. This feature w as proposed as e vidence of loopneb ular interaction and conrmed by decameter ( Ode gard 1986 ) and dif fuse X-ray emission ( Leahy et al. 1986 ) observ ations. Later high ener gy (100 MeV) g -ray images from EGRET ( Jaf fe et al. 1997 ), re v ealed a feature partly coincident with the laments and apparently signicant (7 s ) o v er e xpected dif fuse emission. If real, these g ener gy photons w ould be a product of the interaction of char ged particles with the dense ambient medium at the shock re gion. Recently the HEGRA system of atmospheric Cerenk o v telescopes at IA C w as used to calculate the cosmic ray emission from the loop-neb ula interaction re gion, b ut no signicant T eV ener gies were found ( Aharonian et al. 2004 ). When the Rosette Molecular Cloud w as conrmed as a re gion of star formation, Gre gorio-Hetem et al. ( 1998 ) used R OSA T data again, this time to map the MonR2 cluster and the Rosette Molecular Cloud areas in order to conrm a correlation between star forming cores and clusters of X-rays sources. The y found strong X-ray emission in NGC 2244, the ridge of the cloud (A1-2 in g 2–4 ), and at the cloud core area (A1-1), b ut the resolution w as poor and indi vidual sources could not be resolv ed. The y suggested that molecular cores kno wn to ha v e acti v e star formation b ut f ailing to sho w signicant X-ray emission, could be predominatly forming lo w-mass stars. The y also suggested that detectable X-ray counterparts are in most cases Herbig/AeBe or T T auri stars, as found in NGC 2244 ( Li et al. 2002 ).

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22 The more recent observ ations of the Rosette Comple x done with Chandra ( T o wnsle y et al. 2003 ) ha v e resolutions of only a fe w arcseconds, thus allo wing for the detection of X-ray counterparts for 75% of the OB members of NGC 2244. One of the most interesting results of this study w as the conrmation of a second, soft dif fuse emission which probably originates from the O star winds and is later brought to thermalization by wind-wind interactions or by the shock with the surroundings, in this case the molecular cloud (see gure 2–7 ). This X-ray plasma surrounds the OB association and lls the neb ula ca vity completely 6:35:00 30 34:00 30 33:00 30 32:00 31:30 5:00:0055:0050:0045:0040:0035:0030:0025:0020:00 05:00 Right Ascension (J2000)Declination (J2000) Figure 2–7: A 0.5-2 k eV Chandra image of the Rosette Comple x. The emission has been smoothed to highlight the soft dif fuse emission that originates in the nebula and propagates into the molecular cloud. Credit: T o wnsle y et al. and N ASA/Chandra X-Ray Observ atory (2003).

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23 2.4 The Rosette Molecular Cloud: Embedded P opulations The coincidence of massi v e clumps and luminous IRAS sources pointed out by the study of W illiams et al. ( 1995 ) strongly suggested that star formation had already tak en place across the molecular cloud. Ho we v er the poor spatial resolution of the IRAS point source surv e y did not allo w the resolution of indi vidual members of an embedded population. Early near -infrared studies (e.g. Perez et al. 1987 ) did not co v er the molecular cloud areas, and optical photometric studies were incapable of detecting obscured populations. An e xploratory near -infrared surv e y (JHK) by Phelps & Lada ( 1997 ) that made use of the imager SQIID nally conrmed the e xistence of embedded clusters in some of the most massi v e clumps from the list of W illiams et al. ( 1995 ) that were associated with an IRAS source. The y were able to distinguish se v en deeply embedded clusters with bright neb ulosities, and suggested that clumps not forming a cluster might not be physically bound. The location of the se v en Phelps & Lada ( 1997 ) clusters is sho wn in Figure 2–8 Complete area co v erage of the Rosette Comple x in the near -infrared w as rst accomplished with the release of the All-Sk y 2MASS surv e y catalogs. The 2MASS surv e y w as a major gain in data uniformity b ut unfortunately not in sensiti vity Due to the distance to the Rosette (1.6 kpc), the 2MASS completeness limit (K=14.3 mag) is not deep enough to study the lo w mass end of the IMF So, what is the ne xt logical step in the study of the Rosette Comple x? The e xistence of embedded clusters in the Molecular Cloud means that the cloud is acti v ely forming stars and that at least a fraction of the ne w stars were formed in clusters from the collapse of some of the most massi v e clumps of molecular gas. This leads to tw o problems of importance:2.4.1 Dominance of Cluster F ormation in the Rosette Complex The rst problem is to determine if star formation in the RMC leads to a dominance of rich clusters. Are there an y lo w density groups as well? Is there an y e vidence for a

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24 Figure 2–8: The location of the clusters identied in the study of Phelps & Lada ( 1997 ). The background image is an optical plate from the Digital Sk y Surv e y The contours indicate 12 C O emission from the maps of Blitz & Stark ( 1986 ) distrib uted population? F or e xample, Carpenter ( 2000 ) sho wed that the MonR2 re gion, might be harboring a lo w density population, counting for up to 9% of the total number of young stars. The nature of such lo w spatial density members is not clear as it could be formed independently of the cluster population, b ut also could be the result of the dispersal of an older high spatial density population. Another study that attempts to account for the contrib ution of distrib uted populations in star forming clouds w as done by Li et al. ( 1997 ), who found that for the L1630 cloud, where Lada et al. ( 1991b ) found unequi v ocal dominance of cluster formation. The fraction of infrared e xcess stars in the inter -cluster areas of the cloud w as found to be v ery small, suggesting that the contrib ution of lo w-density formation w as almost ne gligible.

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25 W e need to consider that the RMC is located 4 times further a w ay than the well studied Orion or Perseus molecular clouds ( d = 300 500 pc), where reasonably deep observ ations can easily detect lo w mass stars. Equi v alent detections in the RMC w ould need observ ations at least 3 magnitudes deeper Furthermore, the RMC is located in the direction of the galactic anticenter (l=207de g ) and closer to the galactic disk (b=-2 ) than Orion (b=-16.3) or Perseus (b=-20.6). As a result, the density of the eld population to w ards the RC is v ery high, and an y f aint, lo w mass stars are probably well mix ed with fore ground and background sources. These problems w ould mak e v ery dicult to detect other clusters or a lo w density population by means of single band stellar density counts, as in other cloud surv e ys ( Lada et al. 1991b ; Carpenter 2000 ).Also, depth limited databases lik e 2MASS are not sensiti v e enough to study the Rosette Comple x. Deep multi-w a v elength photometry capable of rendering infrared colors, for e v en lo w mass populations is necessary to separate members from eld stars by means of accurate e xtinction statistics and counts of infrared e xcess stars. 2.4.2 The Hypothesis of Sequential Star F ormation The second problem to be understood, relates to the physical processes that led to the formation of stellar clusters in the Rosette: Are those processes similar to those occurring in other star forming clouds?. The current hypothesis is that the formation of star clusters in the RMC w as possibly stimulated by the interaction of the HII re gion and the cloud. The e xpansion of the Neb ula via the ionization front generated by the strong stellar winds of the massi v e association NGC 2244, results in a shock front which interacts with the gas of the molecular cloud, as sho wn in some of the studies mentioned abo v e. The hypothesis is that the shock front directly stimulated the collapse of clumps which then formed the clusters. This model is kno wn as sequential star formation (SSF), and w as de v eloped theoretically by ( Elme green & Lada 1977 ). In the study of W illiams et al. ( 1995 ), it w as sho wn that the cloud had lar ger v alues of e xcitation temperature, clump density and possibly star formation ef cenc y near the

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26 HII re gion. Ho we v er there are not signicant dif ferences among characteristics of cluster forming clumps, namely mass, size or line width across the cloud. Could this mean that other massi v e clumps, either those not included in the areas of the Phelps & Lada ( 1997 ) surv e y or those not associated with a luminous IRAS point source could also ha v e formed stars recently e v en if their location is not f a v orable with respect to the shock front? In other w ords, ho w feasible is the hypothesis of SSF? The detection of additional embedded populations w ould allo w us to determine for once if star formation is preferentially located near the shock front of the neb ula e xpansion. W e might also be able to nd a relation between the characteristics of the embedded clusters and their distance to the HII re gion that could support or discard the hypothesis of SSF

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CHAPTER 3 A NEAR-IR SUR VEY OF THE R OSETTE COMPLEX: OBSER V A TIONS 3.1 The FLAMINGOS GMC Sur v ey As we discussed in Chapter 1 a systematic and thorough in v estigation of the young star population of GMCs is the k e y to understanding the global aspects of the problem of Star F ormation. Historically the lar ge distances and lar ge angular sizes of GMCs, made it dif cult and costly to perform surv e ys of embedded populations which could render both photometric depth and area co v erage. T echnological limitations were also a f actor with infrared detectors ha ving v ery small areas: until the early 1990s, near -infrared arrays were no lar ger than 256 256 pix els which resulted in rather poor resolution and a small eld of vie w (FO V). F or e xample, the surv e y of the re gion L1630 in the Orion Molecular Cloud by Lada et al. ( 1991b ) used a 58 62 pix el de vice which rendered a FO V of only 1 0 1 0 and thus required of 2800 images to co v er an area of approx. 0.7 square de grees in the K band. Lar ge de vices were de v eloped then, with the instrument SQIID ( Ellis et al. 1993 ) being the rst v ersatile instrument to allo w a high resolution and a lar ge FO V (1024 1024 InSb de vice with simultaneous quadrant detection in J,H and K). This camera w as used for the rst time to surv e y lar ge areas of star formation re gions in multi-band mode, lik e rho Ophiuchi ( Barson y et al. 1997 ) and the Rosette Molecular Cloud (PL97). Near the end of the decade, the rst 2048 2048 HgCdT e de vices for the use in astronomical instrumentation were de v eloped ( K ozlo wski et al. 1998 ), opening e v en better possibilities. The instrument FLAMINGOS ( Elston 1998 ), de v eloped at the Uni v ersity of Florida, tak es adv antage of the 4 million pix el detectors by being designed as a combination wide eld near -IR imager and multi-object spectrometer The camera has a L yot stop wheel with a number of stops customized to recei v e dif ferent 27

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28 input beams slo wer than f/7 and therefore can pro vide a wide range of pix el scales for imaging. F or e xample, on the Kitt Peak 2.1m telescope it renders 0.606 00 pix els and a 20 0 20 0 arcminute FO V This particular setup mak es FLAMINGOS an e xcellent surv e y imager as entire square de gree areas can be surv e yed with a fe w observ ed elds. The instrument has a suite of four lters: J, H, K and K s which co v er the whole near -IR w a v elength range from 1.6 to 2.2 m. Intended as one of the rst lar ge scale applications of the instrument, the NO A O surv e y program T owar d a Complete Near -Infr ar ed Spectr oscopic and Ima ging Surve y of Giant Molecular Clouds (PI E. A. Lada) is dedicated to the global study of se v eral giant molecular clouds using FLAMINGOS. One of the tw o main goals of the surv e y is to do a complete imaging co v erage in J, H and K of comprehensi v e areas of the clouds with a photometric depth that assures co v erage do wn to the Hydrogen Burning Limit (HBL). F our important GMCs were selected for this surv e y: Orion B, Perseus, Monoceros, Cepheus, Serpens and the Rosette. Observ ations for the surv e y program ha v e been carried out o v er the course of 6 winter observing seasons from 2000 to 2005. The surv e y w as done at the 2.1 and 4.0m telescopes of the Kitt Peak National Observ atory where FLAMINGOS is a comissioned instrument. Although FLAMINGOS is suitable for multi-object spectroscop y (MOS) and imaging mode in both telescopes, we preferentially performed imaging at the 2.1m telescope, where FLAMINGOS has a lar ger FO V while the 4.0m telescope has been used essentially for MOS. The imaging observ ations for the GMCs tar gets were carried out iterati v ely with Orion B and the Rosette being the rst clouds to be completed. After reduction of the rst batch of observ ations, a rst quality assessment w as performed by members of the team and collaborators at the Center for Astrophysics in Cambridge, Massachussetts during the f all of 2003. Those elds that yielded poor results were assigned for re-observ ation during the winters of 2003 and 2004.

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29 In the particular case of the Rosette Comple x, a total of 22 FLAMINGOS elds were observ ed during the winters of 2001 to 2004 in the 3 a v ailable lters J, H and K, resulting in a v ery complete co v erage of the Rosette Neb ula the Rosette Molecular Cloud areas. W e selected the area of the Comple x to be co v ered from the 12 CO and 13 CO emission maps of Blitz & Stark ( 1986 ) and the 25 m emission map from the IRAS surv e y T wenty of our 22 elds are adjacent, while 3 of them (areas 4, G1 and G2 ) were added to enhance the quality of the observ ations in some particularly interesting re gions. In order to account for the eld contamination, tw o control elds were observ ed at close distance from the surv e y areas b ut a w ay from the main molecular cloud emission. The control elds were observ ed with an equi v alent method to the main surv e y elds, and ha v e the same depth as an y of our on-source elds. A map sho wing the positions of the observ ed elds in the conte xt of the molecular gas emission ( 12 CO) and the 25 micron IRAS ux in the area can be seen in Fig 3–1 F or each eld we aimed for a total of 1000 sec. on source inte grations in each lter (for some elds, weather conditions and defecti v e frames k ept us a tad belo w this goal), which w as done by obtaining a number of short, dithered e xposures. F or the J and H lters, we used dither e xposures of 60 sec. each each dither and for the K band, with a higher sensiti vity we used 20 or 30 sec. dithers, depending on the weather conditions. Details of the observ ations, including dates, total inte gration times, a v erage seeing, and airmass for each eld can be consulted in Appendix A 3.2 Data Reduction 3.2.1 The Data Reduction Pipeline: LongLegs Each FLAMINGOS indi vidual image is stored as a FITS le with a size of 16 Me gabytes. Each eld is observ ed in three lters and requires combining groups of dithered pointings. The resultant amount of data for the surv e y is therefore v ery lar ge, and required the de v elopment of automated processing pipelines for reduction and photometry

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30 Figure 3–1: Scheme of the Uni v ersity of Florida/NO A O Rosette Molecular Cloud Surv e y The box es delimit indi vidual FLAMINGOS elds (20 20 0 after trimming) o v er an image of the IRAS 25 m emission in the re gion. The labels at the left side of each box refer hereafter to the elds detailed in Appendix A and the te xt. Light solid contours represent the e xtension of the Rosette Molecular Cloud in CO emission from the surv e y of Blitz & Stark ( 1986 ). Crosses mark the centers of kno wn embedded clusters from the pre vious study of Phelps & Lada ( 1997 ). Our image reduction pipeline, nicknamed LongLe gs and programmed by the author is a standard IRAF routine script di vided into three main phases:

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31 During the rst phase, the pipeline rejects defecti v e images and remo v es bad pix els. Then, it applies a 3rd de gree polynomial linearization correction for e v ery image on a pix el by pix el basis (IRAF routine IRLINCOR). Dark and at eld data groups are combined into master at and dark elds, which are then used to create bad pix el masks. The second phase of LongLe gs is a tw o-step preparation of pre-combined data. The algorithm combines groups of 8 adjacent images to create a local sk y Then, after each data frame has been sk y-subtracted and di vided by the normalized at eld, the program reconstructs the indi vidual images dither pattern based on the positions of the 200 brightest sources in each frame. The program does a rst combination of data, e xtracts the positions of sources with ux es lar ger than a pre-selected sigma le v el, and masks them out from indi vidual images to create a ne w set of ”starless” local sk y frames. These are used for a second pass of sk y subtraction, at eld di vision and shift-and-add combination. The nal result of this phase is a set of precombined frames and a rst combined image with analysis quality that only lacks a geometric distortion correction. On the third phase the pipeline program corrects for geometric distortion, using a sixth order Chebyshe v polynomial solution map constructed from the positional distortions of a 20x20 pinhole grid mask that is pre-imaged each time the instrument is corrected internally (corrections indicated slight v ariations in the geometric distortion from season to season). The pre-combined data is also re-sampled to half-size pix els, the dithers are centroid corrected, and re-combined into a nal image that is 4096 4096 pix els in size and is ready for the photometry pipeline. 3.2.2 The Photometry and Astr ometry Pipeline: PinkP ack Our photometry and astrometry pipeline, nicknamed PinkP ac k and programmed by Joanna Le vine ( Le vine 2006 ), performs stellar prole tting (also kno wn as Point Spread Function or PSF tting) photometry on a LongLe gs nal product. The script also uses standard IRAF-D A OPHO T tasks ( Stetson 1987 ), e xcept for the detection, which is performed using the S-e xtractor algorithm ( Bertin & Arnouts 1996 ). A full

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32 description of Pinkpack can be found in Le vine s PhD thesis. W e will only mention that the pipeline gi v es out a full photometric calibration and an astrometric solution with respect to the 2MASS All Source Catalog Release data. Calibration of data is done in the range K = 11 to 14.5 mag. Once a photometric catalog is obtained and an astrometric solution is calculated, this pipeline combines the data from dif ferent lters into a nal mer ged catalog that contains, for each object, an ID, nal RA-DEC coordinates, pix el positions (in the K band image), and photometry for all the bands, including prole tting uncertainties. The photometry pipeline also has the option of creating and subtracting a median v alue image for a specic eld in order to enhance the detection of sources in re gions with bright neb ulosities. W e used this option in all frames that contained bright neb ulosities, although doing the same in non-neb ulous re gions had no ef fect whatsoe v er on the number of detections obtained. 3.3 Completeness of Sample In order to estimate the completeness of our sample, we performed intensi v e articial star e xperiments in 3 selected elds of the surv e y each one considered to be characteristic of a type of re gion: cro wded (with lo w e xtinction), sparse (high e xtinction, no neb ulosity) and with bright neb ulosity Our main goal w as to determine mean v alues of completeness to apply to the entire surv e y for our statistical purposes. The completeness limits for the re gion containing bright neb ulosity emission were not af fected in a greater w ay than in zones of higher stellar density although in both cases the e xperiments performed slightly better in the sparse elds (see Figures 3–2 and 3–3 ). F or all the elds, the articial stars were added partially in consecuti v e annuli of 250 pix els from the center of the frame. F or each annuli, 100 articial images with 50 stars in an uniform distrib ution were created based on the resultant PSF proles from Pinkpack for that specic eld. Their magnitudes were adjusted according to the mean zero point v alue calculated from the 2MASS calibration. The resultant images were then reduced

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33 Figure 3–2: Results of the articial star e xperiments described in section 2.4. a) the left side panels sho w the turnof f magnitudes (limits of 90% object reco v ery) by lter The a v erage v alues in this graphs were used as our general completeness limits for the surv e y with the same set of parameters as the original frame, and the positions and magnitudes of the articial stars were reco v ered using the XYXYMA TCH routine from IRAF The stars in the reco v ery catalogs were di vided by brightness in bins of 0.25 mag, and the completeness limit w as calculated as the bin at which the reco v ery fraction descended belo w 90 percent.

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34 Figure 3–3: Results of the articial star e xperiments separated by type of eld. W e detected a subtle v ariation of the reco v ery limits in the case of too cro wded or too neb ulous elds. The resultant a v erage completeness limits are K=17.25, H=18.00 and J=18.50 mag within the limits of acceptable focus quality of the images (see section 3.4 ); these results rapidly de grade in the areas of high optical distortion. Ho we v er these stand for no w as some of the deepest observ ations of the re gion, going about 3 magnitudes f ainter than 2MASS, and thus assuring the detection of stars around and belo w the HBL.

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35 3.4 P ositional Corr ection of Photometry During the assesment of data quality it w as noticed that our pipelines had dif culties adjusting correctly the PSF in certain areas of the chip. The problem w as w orse to w ards the corners of the images, where the stellar proles were in some cases clearly aberrated and e v en presented prominent comas. It is kno wn that the parabolic shame of the primary mirror has an ef fect on lar ge detectors, which can be usually corrected with a second de gree surf ace v ariation of the PSF prole, b ut apparently the distortions we observ ed had a dif ferent origin, because the distortions are not symmetrical, i.e., the four corners of the images are af fected dif ferently The distortions can also w orsen with poor focusing and bad weather (i.e. mediocre seeing v alues). One hypothesis based on optical path simulations (Eik enberry S. Uni v of Florida, personal communication) is that the alignment between the primary and secondary mirrors of the KPNO-2.1m telescope has lost accurac y along the years, af fecting the symmetry of the focus and shifting the center of optical alignment. This def fect is unfortunately enhanced by the lar ge FO V of FLAMINGOS. The size of the af fected area v aried slightly from season to season. The area with minimal distortion is nearly circular with a center that f alls systematically on the pix el position (3000,2400) for observ ations made before the f all 2004, and on the position (3200,2170) for more recent observ ations. The radius of this area within which the PSF v alues stay uniform is v ariable, with an a v erage of 3200 600 pix els depending on the observing conditions, mainly seeing v alue, which is af fected respecti v ely by the airmass and the weather conditions at the observ atory F or most of our elds, about 75-95% of the area of the detector contained minimal distortion, with reasonably smooth PSF proles and small ( < 0 : 025 mag) photometric dif ferences with respect to the 2MASS catalogs. Outside this area, the optical distortion increases quickly and therefore the shape of the stars and the PSF proles de graded to the point that stars presented noticeable aberration comas and lar ger PSF FWHM

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36 v alues,especially at the tw o eastw ard corners of the detector which resulted in poor ttings to the a v erage PSF prole generated from good quality stars in the acceptable area, and generated a net ux loss with respect to 2MASS that raises sharply to 0.5 mag in the bad psf tting areas, independently of the lter In terms of the net output to our photometric catalogs, this ef fect resulted in a v ariation of the photometric calibration zero point (ZPT) v alue across the images, and this af fected the uniformity of the surv e y from eld to eld. In order to correct for this ef fect, we applied a 6th order Le gendre polynomial correction of the ZPT v alues as a function of the radial position with respect to the pix el centers of the optical distortion circles. The solution w as de v eloped and constructed as interacti v e softw are by Andrea Stolte. The correction is applied done by tting a polynomial to a ducial line made by the median v alues of the 2MASS vs FLAMINGOS dif ferences in radial bins of 300 pix els from the optical distortion center The correction w as calculated within the ranges 10.0 to 14.0, 10.0 to 15.0 and 10.0 to 16.0 mag in K, H and J only due to the limitations of sensiti vity of the 2MASS catalogs, b ut w as applied to e v ery star detected by our pipelines. This method allo wed us to reduce the scatter in the ZPT v alues, and to determine (by eld and by lter) which w as the cutof f radius from the minimal distortion center at which ZPT dif ferences with respect to 2MASS rose abo v e a maximum tolerance of 0.3 mag. Inside the area mark ed by this cutof f radius, the polynomial correction reduced the ZPT dif ferences signicantly and this also results in a decrease of the noise in the color terms. The areas located be yond the cutof f circles, to w ards the east (left) corners of the detector ha v e too lar ge optical distortions, and so objects detected in those areas were remo v ed from our nal catalogs Figure 3–4 sho ws schematically the positions and e xtensions of e v ery eld observ ed, as well as the cutof f radii of the ZPT correction for each lter The K band circle, being the most conserv ati v e, al w ays denes the area of the eld that w as k ept for analysis. This

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37 of course, has an e xception for those areas that ha v e an o v erlap with the good quality re gions of another coincident eld, in which case our catalog joining program selected systematically the star from the good frame into the nal catalog. Figure 3–4: The e xtension of the areas of acceptable optical distortion are mark ed for each eld as circles with radii equal to the center of the maximum bin at which the ZPT polynomial correction to the zero points (see te xt and gure 5) can be applied within the detector This ef fect v aries by eld (size of the acceptable area) and lter: the solid, dotted and dashed linestyle circles represent the tolerance radii for J, H and K respecti v ely

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38 In the v arious panels of Figures 3–5 to 3–9 we sho w as an e xample, the ef fects of the polynomial ZPT correction in the area 01 of our surv e y (which coincides approximately with the center of the Rosette Neb ula). The rst group of plots (Figures 3–5 ) sho ws the polynomial ZPT correction applied to the photometric dif ferences FLAMINGOS vs 2MASS as a function of radial position from the minimal distortion center (3000,2400). As can be noticed, the scatter in the zero point per magnitude is clearly corrected and the dif ferences at lar ge radii no w con v er ge closer to zero and stay within a 0.1 mag range. Figure 3–5: Example of the results of the ZPT polynomial correction of the zero point as a function of pix el radial distance from the center of lo w optical distortion (3000,2400) in the FLAMINGOS detector for re gion 1 of our Rosette surv e y in lters J and H. The dots represent matches of the FLAMINGOS data with 2MASS sources in the ranges 11.0 to 15.0 and 11.0 to 16.0 H and J respecti v ely The solid line represents a 6th order Le gendre polynomial t to the median v alues of the scatter in bins of 300 pix els. The dashed lines indicate le v els of 0.1 mag of scatter

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39 Figure 3–6: Same as Figure 3–5 b ut for the K lter In the second group of plots (Figure 3–7 ) we can see ho w the v alues of the photometric dif ferences with respect to the observ ed colors are also reduced intrinsically decreasing the o v erall uncertainty of our photometry The third set of plots (Figure 3–8 ) sho ws the net ef fect of the ZPT correction on the color -magnitude and color -color diagrams. As it can be noticed, the color -magnitude sequences for background and for members in the eld get more conned and better separated, which in the color -color space results in a reduction of the scatter around the zero age and giant sequences. This, consequently reduces the number of spurious detections in the infrared e xcess re gion of the color -color diagram,especially in the area located closer to the intersection of the T -tauri reddening band and the main sequence. In the fourth set of plots of Figure 3–9 we sho w the scatter of these photometric dif ferences for eld 01 before and after the correction is applied. In the left panels of the

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40 Figure 3–7: Example of the results of the ZPT correction in the J H and H K color dif ferences of FLAMINGOS vs 2MASS as a function of magnitude in area 1 of our Rosette surv e y T op and bottom panels indicate color dif ferences before and after the polynomial correction. The color dif ferences are also reduced. The solid line indicates the zero le v el. gure we plot the FLAMINGOS 2MASS dif ferences as a function of magnitude, which has lar ger scatter v alues to w ards f ainter magnitudes. The mean v alue of the dif ferences is closer to zero after the correction. The right panels sho w ho w the correction results also in a reduction of the net photometric scatter measured by the standard de viation of the dif ferences within consecuti v e magnitude bins. The median v alue of these de viations is indicated with a dashed line. 3.5 Quality and Unif ormity of the Sur v ey The o v erall uniformity in the quality of the photometry of our surv e y can be simply assessed by comparing the mean v alues of the scatter in the FLAMINGOS vs.2MASS dif ferences after applying the positional correction, as described abo v e. The median

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41 Figure 3–8: Example of the results of the zero point correction in the color -magnitude and color -color spaces for the sources in area 1. The top and bottom panels represent data and after before the polynomial correction is applied. scatter v alues were compiled for e v ery eld and lter and we use these numbers the indicators of the net photometric quality of an observ ation. The medin scatter v alues for each eld are sho wn in T able A–2 of Appendix A The correction allo wed us to reduce our internal photometric scatter in indi vidual bands by an a v erage of 0.02 mag within the lo w optical distortion areas. The a v erage of these scatter v alues o v er the whole surv e y are 0 : 058 0 : 012, 0 : 064 0 : 018 and 0 : 056 0 : 014 in J, H and K respecti v ely In addition to our completeness limits, which are statistical, there is another set of limits which represent the sensiti vity of the surv e y i.e. the f aintest magnitude at which an object in our catalog can be consider to ha v e good photometric quality within our errors These sensiti vity limits are also dif ferent for each lter and we estimated them

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42 Figure 3–9: Example of the results of the ZPT correction in the v alues and scatter of the FLAMINGOS vs 2MASS dif ferences as a function of magnitude (J band). T op and bottom panels sho w data before and after the correction respecti v ely The numbers on the top of left panels are median and standard de viations within the t range. F or the right panels, the dashed line and the numbers on the top represents the median of the standard de viation per magnitude bin. Solid and dotted lines mark the 0.0, 0.05 and 0.1 scatter le v els. simply as the points where the ducial curv e represented by the a v erage v alues of the FLAMINGOS photometric errors crosses the standard 10 s limit dened as 0.109 mag le v el in J, H and K. The result is sho wn in g 3–10 ; our 10 s crossing v alues are J=19.4, H=18.4 and K=17.7 mag. These v alues are also a good indication of the limits at which our photometric v alues are consistent in quality and are slightly higher than our estimated completeness limits. In f act, it is possible that within certain areas of the elds,especially near the centers of the lo w optical distortion areas, the completeness of the data is in f act higher than the a v erage v alues calculated from the articial star e xperiments, b ut as stated in the section 3.7

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43 Figure 3–10: The solid (ducial) line joins the points which mark the median v alue of the PSF tting photometric uncertainty in a gi v en magnitude bin. Error bars indicate the standard de viation in each bin. The horizontal dashed lines marks the zero le v el and the sensiti vity le v el, estimated at the standard 10s limit of 0.109 mag. V ertical dotted lines indicate the sensiti vity limits, mark ed as the bin at which the ducial curv e generated by the median v alues crosses the 10s line. The dash-dotted line represents the 3s le v el from the ducial line; an y star in our catalog with errors higher than this le v els were rejected from the analysis. this depends greatly on the v ariations of the e xtinction across the Molecular Cloud, and logically e v en the completeness limits can be compromised accordingly in certain spots.

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44 3.6 Construction of Final Catalog 3.6.1 Intrinsic quality: 2MASS Addendum Our photometric uncertainties also de grade relati v ely steeply at the bright end, approximately belo w 11.0 magnitudes in J, H and K. This is because this coincides with the le v el at which the counts per pix el in the detector reach v alues abo v e 3 10 4 counts, right at the limit of linearity and saturation. In order to account for this ef fect, we rejected, for our analysis, e v ery source in our catalog with a magnitude brighter than H=11.0, as this lter seem to be af fected slightly w orse by the saturation ef fect. T o complete this end of the magnitude spectrum, we added 2MASS sources to complete our catalogs within the range 5.0 to 11.0 mag in all lters. The total addendum of 2MASS objects to our surv e y is 798 objects. 3.6.2 Sur v ey Ar ea Mer ging A nal catalog that includes the identications, astrometry photometry and basic information about v ariability in o v erlap areas for all the sources accepted from our surv e y w as created using a catalog mer ging code that put together catalogs for indi vidual elds and calculated the weighted a v erage v alues for matches (duplicates) in o v erlap areas. In the case where a match source w as located in an area of high optical distortion, null v alues were ignored. Most of the re gions in the surv e y are adjacent, so that the o v erlapping re gions were usually small; in f act, the outermost 1 arcmin ribbon of the nal combined images w as in most cases trimmed out of the nal images, because Pinkack only runs on the re gions of the image that contained a combination of at least 50 percent of the indi vidual dithers. Ho we v er re gions 4, G1 and G2 ha v e lar ge o v erlapping re gions within the RMC area and in those cases the selecti v e matching came in handy After mer ging all of the indi vidual catalogs, including the 798 2MASS sources for the bright end, our nal surv e y catalog contained a total of 153,266 objects. F or our analysis ho we v er we restricted our study to those stars with an uncertainty v alues belo w 3.0 s abo v e the median (see Figure 3–10 ). The total number of sources in

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45 this selection is 146,868. Completeness limits at this le v el wer e not af fected as 97.3% of these 3-sigma rejections comprised magnitude ranges abo v e J=18.50, H=18.25 and K=17.50 mag respecti v ely 3.7 Intrinsic Detection Constraints W e consider important to mention that surv e ys with an emphasis on embedded young populations, e v en with the aid of infrared detectors, cannot assure full detection of all embedded cluster members. Most of young clusters are embedded in the remnants of their original molecular gas cores, which are e xpected to carry massi v e amounts of dust and dense gas, to the point that some objects (specially lo w mass protostars) will be so highly e xtinguished that the y will not be detected e v en in the near infrared, and thus a fraction of the true number of cluster stars will be left of f the countings. In addition, cluster populations will al w ays be mix ed with a signicant number of fore ground and background objects, whose contrib ution has to be estimated from the control elds. Contrib utions of eld stars are corrected by e xtinction ef fects, counted per magnitude bin and subtracted from equi v alent counts on-eld. After these subtractions are applied, nal countings for the number of members in a cluster re gion are only statistical estimates and cannot determine the membership of indi vidual sources. Another related f actor that has to be tak en into consideration is that lo w mass premain sequence stars become f ainter at older ages, with the consequence that at a certain sensiti vity limit the y will just not be detectable ( Carpenter 2000 ). This sensiti vity limit depends on the age and mass distrib ution, as well as e xtinction, and the net ef fect is only kno wn accurately for a fe w re gions. In our gure 3–11 which is similar to gure 18 from ( Carpenter 2000 ), we try to sho w the intrinsic limitations on mass and age detection that our RMC surv e y has. F or our completeness limit of K=17.25 we e xpect to detect stars well be yond the HBL if all the stars were younger than 2.5 Myr and no e xtinction w as present in the line of sight; if stars are older the range of observ able stellar mass is reduced. In a typical molecular

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46 Figure 3–11: Contours of equal K magnitude v alue as a function of stellar mass and age at the estimated distance of the Rosette Molecular Cloud (1600 pc, distance modulus=11.02). F or the construction of these plots we used the pre-main sequence e v olution models of D'Antona & Mazzitelli ( 1997 ). The solid lines represent iso-magnitude le v els with no e xtinction, while dotted lines represent the same le v els with a A V =5.0 mag e xtinction. The contour at K=17.25 coincides with the completeness limit of our FLAMINGOS surv e y and the dotted v ertical line marks the 0.08 M HBL limit. cloud, typical e xtinction v alues range from a fe w to 50 magnitudes in visual w a v elengths, which w ould cause the completeness limits to compromise up to 5 magnitudes in K, causing the range of detectable ages and masses to be e v en shorter in some areas.

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47 Most of the kno wn clusters in the Rosette are deeply or partially embedded in their parental cores (see 5 ), which means that their ages could be possibly no older than 1-2 Myr (the most recent spectroscopically estimated age of NGC 2244 is about 2 Myr), and thus, more than 80% of the embedded stars could be detected, depending on the e xtinction le v el, b ut this ef fect is not uniform e v en at sub-cluster scales.

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CHAPTER 4 NEAR-IR SUR VEY : AN AL YSIS AND RESUL TS 4.1 Intr oduction As discussed in Chapter 1 the study of embedded clusters in Giant Molecular Clouds (GMCs) is of capital importance to understand the problem of star formation: at the embedded stage, star clusters ha v e not e v olv ed signicantly and therefore their densities and mass distrib utions are still close to the original fragmentation of their forming cores. From observ ations of embedded clusters in nearby GMCs ( d < 2 kpc), there is good e vidence that a major fraction of the stars are formed in a cluster en vironment, with ric h clusters (100 members or more) clearly dominating o v er small groups, as the y contain more than 80% of the embedded stellar population in GMCs. Unfortunately a v ailable catalogs of embedded clusters are still incomplete at the small cluster sizes. Systematic surv e ys of molecular clouds that are focused on the detection of young clusters are usually limited to those re gions with signposts of star formation (lik e the presence of luminous IRAS sources), which lead to the disco v ery of the richest clusters. Only in a v ery fe w cases has there been an specic search for lo w density groups or distrib uted embedded populations, gi v en that the y are logically more dif cult to detect, especially if the y are mostly composed of lo w mass stars, with lar ge spatial distrib utions and projected against a high background of reddened sources. The RMC is a particularly acti v e star formation re gion, with a lar ge OB association, NGC 2244, whose winds ha v e generated an e xpanding HII re gion. One unsolv ed problem is to determine if the shock front generated by this photodissociation b ubble w as the principal trigger of the formation of the observ ed embedded clusters found by ( Phelps & Lada 1997 ) in the adjacent, highly structured Rosette Molecular Cloud (RMC) ( W illiams et al. 1995 ). 48

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49 All of the PL97 clusters are associated with a luminous IRAS source and a massi v e molecular clump, and until no w no other study has been able to determine the e xistence of additional clusters. One of the main problems is the lar ge distance to the Rosette: the cloud is located at d = 1 : 6 kpc, roughly 4 times further a w ay than nearby clouds lik e Orion or Perseus and for this reason pre vious studies with limited photometric sensiti vity (depth) were unsuccessful at gi ving an y ne w information on the distrib ution of young populations in the Rosette. F or e xample, the 2MASS surv e y is complete only to K = 14 : 3 mag, which is not enough to detect lo w mass stars or deeply embedded high mass stars. This mak es dif cult to detect highly embedded clusters, specially if their local surf ace densities are lo w Another dif culty is that the Rosette is located at a v ery lo w galactic latitude ( b = 1 : 8 2 : 0), which implies a v ery high density of eld objects, and thus, if a search for clusters w as performed using monochromatic w a v elength counts lik e it w as done for other clouds, then corrections for background contamination w ould be lar ge and dif cult to apply Our surv e y is designed to address these problems by a) doing deep observ ations of the re gion, capable of detecting stars close or belo w the HBL and thus impro ving the detectability of pre vious studies, b) replacing the use of monotonic w a v elength counts by emulating the technique of Li et al. in which the detection of young populations is done by photometric color selection, and c) using the method of Nearest Neighbors ( Casertano & Hut 1985 ) to distinguish areas with surf ace densities intrinsically lar ger than the o v erpopulated eld. Among the main goals of our surv e y are to study the characteristics of the kno wn clusters, to determine if there are more, and to study their distrib ution across the comple x, and in the conte xt of NGC 2244. Because our observ ations are deep enough to detect lo w mass stars we should be able to trace well the structure and e xtension of embedded populations in the Rosette Comple x, adding v aluable information about the nature of stellar nurseries in GMCs.

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50 4.2 Analysis 4.2.1 The Near est Neighbor Method Single band monotonic detection of embedded clusters is done by subtracting normalized control eld counts (corrected by e xtinction) from counts to w ards the cloud line of sight. This is e xpected to gi v e the total number of e xpected members of a certain re gion, and clusters are dened as re gions with surf ace densities signicantly higher than the eld. Unfortunately the method is biased because it will only be able to detect v ery lar ge or v ery dense clusters. Another problem is that the subtraction of eld counts is more dif cult for more distant clouds, because members are f ainter and are mix ed with a lar ger number of fore ground and background stars. The Rosette Comple x is located 4 times further a w ay than other star forming comple x es where systematic searches for embedded clusters ha v e been performed, lik e Orion or Perseus. Furthermore, the Rosette is located at a v ery lo w galactic latitude ( b = 2) and to w ards the anticenter of the Galaxy ( l = 210), which results in an intrinsically high density of eld sources located in the fore ground and background of the cloud. F or e xample, one typical of f-source eld in the Rosette observ ed with FLAMINGOS can ha v e an a v erage of 6 10 3 sources, almost 5 times higher than a eld in Orion. T o detect embedded populations in the Rosette we applied a density selection technique, the Nearest Neighbors Method (hereafter NNM), to distinguish populations with surf ace densities abo v e the uniform eld le v els. This method has already been applied succesfuly to nearby clouds by Ferreira et al. ( 2005 ) and gi v es reasonable results for a lar ge distance re gion lik e the Rosette. Ho we v er we impro v ed the method with the use of a color selection to separate the youngest members in a eld as near infrared e xcess (IRX) objects, assuring the detection of embedded clusters by increasing their probabilities of membership. W e e xpect this combined approach to gi v e a non-biased detection of young stellar groups and to gi v e insight into their nature at the same time.

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51 Before describing our color selection, we will re vie w briey the terminology and concepts from the NNM rele v ant to this paper A more detailed description of the general use of the method for embedded populations can be found at Ferreira et al. ( 2005 ). The calculation of nearest neighbor densities to detect clusters in a cro wded eld w as proposed by Casertano & Hut ( 1985 ). The y proposed to estimate the local surf ace density of objects in a certain eld from the individual relati v e surf ace density of each object. The generalized form of the indi vidual density estimator is: r j = j 1 p D 2 j ; (4.1) where D j is the distance of the star to its j th member This estimator has only one de gree of freedom, the number j of neighbors used to calculate the local density According to Casertano & Hut the lar ger the v alue of j the smaller the uctuations in the local density estimations due to local irre gularities, which is v ery useful to determine e xtension and structure of lar ge systems. Ho we v er j also denes the minimum number of particles in the smallest substructure to be considered, and so j should be small if structures looser or smaller than typical clusters are to be detected. The y sho wed that j = 6 w as the minimum number at which uctuations could be acceptable for populations of the order of 30 to 1000 particles. The NNM also allo ws the denition of a “density center” and a “density weighted radius” which dene cluster centers and cluster cores, respecti v ely The density center is dened as the density weighted a v erage of the star positions in a eld: ~ X d = i ~ X ( i ) r j ( i ) i r j ( i ) : (4.2) And the density or core radius, R cor e is dened as the density weighted a v erage of the distance of each star to ~ X d ; j :

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52 R cor e = i j ~ X ( i ) ~ X d ; j j r j ( i ) i r j ( i ) (4.3) Using the NNM, clusters can be detected as re gions where a v erage indi vidual r j v alues are lar ger than those of a uniform control eld, and if j is small enough, groups of the order of N = 10 1 stars should be detected without bias. If we follo w the denition of N 35 stars as a minimum number to represent a cluster ( Lada & Lada 2003 ; Adams & Myers 2001 ), then an y groups with less members could be considered loose enough as to account for a non-cluster (distrib uted) population. F or e xample, Ferreira et al. ( 2005 ) applied a j = 20 estimator to 2MASS catalogs of nearby ( d < 1 kpc) molecular clouds and conrmed locations and sizes of clusters with radii as small as 0.3 pc and total number of members as lo w as 20 5 stars. 4.2.2 Detection of Embedded P opulations W e calculated the 20th nearest neighbor densities for stars do wn to the completeness limit in our nal RMC catalog and the control elds, e xpecting to be able to distinguish at least NGC 2244 and the se v en clusters from PL97 1 as re gions with densities unequi vocally higher than the eld. The result w as that NGC 2244 is indeed, v ery well traced as a high density re gion, as were the zone of clusters PL04 and PL05 in the core of the cloud. Clusters PL01, PL03, and PL07 were also distinguishable b ut their apparent e xtensions were not more note w orthy than some groups of stars that rose abo v e the 3 s le v el only because the y coincided with “patches” of lo w e xtinction around the main molecular gas emission. Finally clusters PL02 and PL06, the smallest in the cloud, presented densities belo w the 3 s le v el because their small number of members (around 30) resulted in lo wer than a v erage 20th NN densities, and thus are not distinguishable among the noise le v els of the distrib ution. 1 W e use from no w on the nomenclature ”PL01” to ”PL07” to refer to these clusters.

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53 P articularly cluster PL06, which is associated with the B-type proto binary AFGL961, contains lar ge quantities of obscuring material near its center which mak es dif cult the detection of embedded f aint members e v en in carefully constructed near infrared maps (e.g Aspin 1998 ); an y clusters lik e PL06 will be dif cult to detect in a lar ge j scheme because their number of members will be intrinsically small. Gi v en these dif culties, we repeated the NN analysis using only infrared e xcess (IRX) stars, to assure the detection of only the youngest component of the cloud embedded population, to minimize the contrib ution of the eld and to promote the detection of deeply embedded, lo w surf ace density clusters. Infrared e xcess stars are not e xpected to e xist in the control elds as the y located a w ay from the star forming clouds, and so the density of these objects should be al w ays higher for embedded populations. Also, as the fraction of IRX stars in an embedded cluster is lar ger if the cluster is younger and therefore more embedded, so that high e xtinction re gions could actually ha v e higher IRX o v erdensities, impro ving detection. 4.2.3 Infrar ed Excess Stars In the J H vs. H K color -color diagram, IRX stars f all to the right of the reddening band dened by the projection of the Classic T T auri star (CTTS) locus ( Me yer et al. 1997 ), along the direction of the e xtinction v ector ( Cohen et al. 1981 ). The fraction of stars with infrared e xcess emission in a cluster is kno wn to decrease with time as early stellar e v olution leads to the destruction of disks by photospheric UV radiation, b ut for a deeply embedded population with ages of 1 to 2 Myr lik e the one e xpected in the RMC (the OB association NGC 2244 is estimated to ha v e an age of 1.9 Myr and the embedded clusters cannot be older), the circumstellar disk fraction will be signicant enough and IRX stars counted from JHK e xcess will trace well the presence of the most recent episode of formation (see e.g Lada et al. 1996 ; Carpenter et al. 1997 ). F or our study we dene an IRX star as one with colors that place it 0.1 mag (5 times the standard de viation of the H K uncertainty) to the right of the ZAMS and abo v e

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54 J H = 0 : 47 ( H K ) + 0 : 46 which denes the lo wer limit of the Classic T T auri Star (CTTS) locus of Me yer et al. (see Figure 4–1 ). Figure 4–1: The near -infrared color -color space. The thick dark solid lines represent the loci of the zero age main sequence and the giant branch ( Bessell & Brett 1988 ). The thick colored line is the Classic T -T auri locus ( Me yer et al. 1997 ), which in this diagram is e xtended to the right to lar ge H K v alues and abo v e and belo w by its observ ational error The other dashed lines represent e xtinction along the direction of the reddening v ector indicated by the arro w on the left. The shado w re gion indicates where, under our denitions, stars with possible infrared e xcess emission f all. Stars f alling in the re gions labeled as 1 and 2 colors are usually af fected by spurious detections and high photometric color scatter

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55 The rst constraint a v oids contamination from non-IRX stars with lar ge H K uncertainties located close to right edge of the Zero Age Main Sequence Reddening Band (ZAMSRB). The second constraint helps us to a v oid including objects that locate in the re gion belo w the CTTS line. Sources f all in this re gion mainly due to high color scatter (see section 4.2.4 belo w), ho we v er unresolv ed galaxies with lar ge color dispersions ( Labb e et al. 2003 ) or distant background galaxies with highly inclined reddening v ectors ( Heraudeau et al. 1996 ) may ha v e colors that f all in this re gion of the diagram. F or stellar sources, there are cases in which neb ulosity can add a blue J H component to background stars with a lo w e xtinction v ector and push stars to the re gion, as sho wn by Montecarlo simulations of Muench et al. ( 2001 ). 4.2.4 Magnitude Depth Restriction f or IRX stars The combined ef fects of v ariable seeing quality o v er the seasons, and v ariability of the focus quality across the wide detector eld of FLAMINGOS, resulted in a high intrinsic dispersion of color v alues for our sample, which cannot be eliminated with the zero point corrections or the uncertainty restrictions. In gure 4–2 we illustrate this ef fect in a contour le v el color -color diagram all of the stars in our w orking catalog. In the diagram, made with a n yquist box size of 0.1 mag, the lo west le v el sho wn represents the mean color -color space surf ace density and each subsequent le v el represents a step of 1 standard de viation. There is a noticeable bloating in the dispersion of colors at both sides of the ZAMSRB at the core of the diagram, near the re gions of lo west e xtinction. The color dispersion is lar ger for f aint stars. The scatter in H K has a major increment at approximately K=15.75, where the eld object density increases signicantly In Figure 4–3 we sho w equi v alent contour le v el color -color diagrams made with separate samples for stars with K < 15 : 75 mag and for stars with 15 : 75 < K < 17 : 25 mag respecti v ely The diagrams sho w that the scatter for the bright end bins is smaller than for the f aint end bins. Our calculations indicate that the standard de viation of the H K color uncertainties is twice as lar ge for the f aint end bins (0.11 vs. 0.21 mag). This ef fect is

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56 11.0 < K < 17.25 Figure 4–2: Contour le v el color -color diagram for all stars in the FLAMINGOS RMC surv e y within the restrictions described in section 2.1. The diagram w as constructed using a sampling box of size 0.1 mag. The lo west le v el represents the mean v alue of the object counts, at 825 de x 2 Subsequent le v els represent steps of 1 sigma (3550 de x 2 ). slightly w orse for some elds which were observ ed under less f a v orable weather conditions or at a higher than a v erage airmass. There are also re gions of the surv e y where the scatter is smaller and the quality of the colors is k ept to f ainter limits (see section 4.3.2 ). The statistical cuts we present are mostly conserv ati v e and assure the uniformity of our statistics across the entire surv e y area.

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57 a) 11 < K < 15.75 b) 15.75 < K < 17.25 Figure 4–3: Contour le v el color -color diagrams for stars in the FLAMINGOS RMC surv e y di vided in tw o ample groups of brightness. Both diagrams were constructed with a Nyquist box size of 0.1 mag. The diagram a) sho ws the distrib ution of colors for `bright' stars within 5 : 0 < K < 15 : 75 mag, and the diagram b) is for `f aint' stars within 15 : 75 < K < 17 : 25 mag. The contour le v els start at the mean le v el (360 and 466 de x 2 for a) and b) respecti v ely) with subsequent steps of 1 sigma (1770 and 1910 de x 2 for a) and b) respecti v ely). The high scatter in near infrared colors af fects directly the calculation of the number of IRX stars in the surv e y which locate to the right of the ZAMSRB. W e performed Montecarlo e xperiments in which we simulated the colors of stars dra wn from a model population with an age of 1 Myr ( D'Antona & Mazzitelli 1997 ) located at the distance of the Rosette and we added to our simulated stars, color errors and e xtinction similar to those observ ed in the surv e y areas. W e found that the resultant number of stars with colors similar to those of IRX stars w as 5 times lar ger for stars in the f aint end. Because we are basing our analysis in the detection of infrared e xcess sources, we had to limit the primary aspect of our analysis, the identication of embedded clusters, to those stars in the bright end of the sample to assure an IRX sample with a minimum of contamination. Ho we v er these bright IRX sources are only helping us to tr ace the location and rough e xtension of clusters, and although our color uncertainties are high,

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58 indi vidual K band magnitudes are still good within 0.1 mag, which allo ws us later to include stars do wn to the HBL in the areas traced by the IRX sources and calculate correct luminosity functions with a generous bin resolution of 0.25 mag. Also, a photometric depth limit of K=15.75 means a sample still almost 1.5 magnitudes deeper than 2MASS and is equi v alent (for dw arf type stars) to a stellar mass range of 0.09 to 0.18 M for a population of 1 Myr embedded in a cloud with a typical e xtinction of 0 to 10 visual magnitudes ( D'Antona & Mazzitelli 1997 ). Thus, we should be able to count IRX populations slightly abo v e the HBL for a typical young cluster 4.2.5 Near est Neighbor Analysis f or Infrar ed Excess Stars A preliminary e xploration using f alse color image combos of the less populated clusters, PL01, PL02 and PL06, indicated that the typical number of stars in a modest size cluster that can be detected ”by e ye” within areas of bright neb ulosity usually coincident with embedded cluster cores, could be of the order of 30 members. This is close to the minimum number that denes an association of stars to be a cluster and belo w that, an y groups could be considered a distrib uted population. The e xpected JHK infrared e xcess fraction in an embedded cluster less than 3 Myr old is 20 to 60 percent, so the number of IRX sources is much smaller than the number of sources in the full catalogs, and cluster ha v e to be identied with less neighbors. Because of this, instead of j = 20 as in FL06, we selected a v alue of j = 10, which assures the detection of clusters with less than 20 IRX members (60 to 100 total members if the fraction of stars with circumstellar emission is 20 to 60%) b ut gi v es local surf ace densities with 15% more accurac y than the minimum j = 6 described by Casertano & Hut ( 1985 ). Also we are able to determine the e xistence of populations distrib uted in groups almost three times less dense than the minimum e xpected for a cluster W e detected a total of 1168 34 sources IRX sources under our denition in the bright end sample. In Figure 4–4 we sho w their 10th Nearest Neighbor distrib utions of

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59 Figure 4–4: Nearest Neighbor distrib utions for bright IRX stars. The top panel sho ws the distrib ution of 10th neighbor distances. The bottom panel is the distrib ution of 10th neighbor densities. In the top panel line A indicate the limit of distances shorter than 1.0 pc, while the dashed line B indicates the midpoint v alue at 1.83 pc. At the bottom panel the equi v alent limits in density space are also indicated. distances, D 10 and local surf ace densities r 10 The mean v alue found for D 10 w as 1.83 pc which corresponds to a r 10 = 0 : 2 ( 0 ) 2 This limit is indicated in Figure 4–4 F or the control elds we found 3 sources that had IRX colors do wn to a maximum brightness of K = 15 : 75. Ho we v er we added another 16 sources which f all to the right of the reddening band belo w the CTTS line b ut w ould ha v e IRX J H colors if reddened by

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60 an a v erage v alue of A V = 5 : 0 mag, typical of the cloud re gions. The mean 10th neighbor density among these 19 sources w as 0.18 ( 0 ) 2 which compares relati v ely well with the mean v alue 0.2 ( 0 ) 2 of the distrib ution of distances in the surv e y areas, so that we considered a round v alue of 0.2 ( 0 ) 2 as a bac kgr ound limit, belo w which we cannot assure that an IRX source has a density high enough to be distinguished from the eld. The minimum v alue of the D 10 distrib ution in the surv e y to be 0.145 pc, which represents a density of 29.5 ( 0 ) 2 and is a good estimation of the typical local surf ace density in the central re gions of RMC clusters. The midpoint between this minimum distance and the mean is 0.987 (roughly 1 pc) which corresponds to a r 10 v alue of approx. 0.6 ( 0 ) 2 This can be considered as a good estimate of the a v erage embedded cluster size in the RMC. 4.2.6 Identication of Clusters In Figure 4–5 we sho w the location of IRX stars with the le v els of density described abo v e. All of the kno wn clusters seem to be traced well in this selection, and we conrm that the y are the main re gions of star formation in the Rosette Comple x. In Figure 4–6 we present a contour le v el plot of the local surf ace densities calculated with the NN method ( j = 10). The contours were constructed using a Nyquist sampling box of 90 00 W e dene a cluster as a re gion for which a closed contour at 0.2 ( 0 ) 2 contains at least 10 IRX sources. Using this denition we found, in addition to NGC 2244 and the se v en PL97 clusters, 4 additional areas that arise as signicant b ut ha v e not been studied before: The rst is a re gion at the “Core” of the cloud, to the east of cluster PL05 and south of cluster PL04. This re gion, which we designate as RLE08, contains a lar ge number of highly reddened sources, dif fering from the clusters PL04 and PL05 which ha v e a number of sources already visible in DSS plates and can be considered partially emer ged. RLE08 appears to be a more recent episode of formation in this ample zone of formation at the

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61 Figure 4–5: The location of IRX stars in the Rosette surv e y with brightness K < 15 : 75 mag. The ”plus” symbols are IRX stars with 10th neighbor densities higher than 0.2 ( 0 ) 2 while black dots are stars with densities belo w 0.2 ( 0 ) 2 W e also indicate the e xpected position of the center of NGC 2244. Contours indicate le v els of CO emission in steps of 20K km s 1 The dotted line indicates the limits of the surv e y co v erage. center of the cloud, in which clusters PL04 and PL05 are the lar gest and most brilliant clusters. The second is a substantially lar ge, highly reddened cluster located in the southeastern edge of the cloud, designated as RLE09. Along with RLE08, these clusters are clear e xamples of clusters which are located in re gions with lar ge e xtinction v alues and thus

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62 Figure 4–6: Identication of clusters in the Rosette Comple x. The contours indicate 10th Nearest Neighbor densities and were constructed with a n yquist box size of 1.5 arcmin. Labels for indi vidual clusters are e xplained in te xt. The dotted thin lines indicate the 15.0 K km s 1 le v el of 12 CO emission, which we use to dene the e xtension of the main molecular cloud re gions. ha v e a v erage surf ace densities comparable or lo wer than the eld. Ho we v er the y contain a lar ge number of red sources that are easily distinguishable in JHK f alse color composite images and such a lar ge number of IRX sources that the y stand out clearly as embedded clusters.

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63 A third ne w cluster is located to the east of NGC 2244, in the re gion of the cloud identied as NGC 2237, which is well kno wn for its content of gas pillar structures ( Carlqvist et al. 1998 ), and thus we assign it to that name. The e xistence of this cluster w as also suggested by Li ( 2005 ) in their study of 2MASS data. NGC 2237 is distinguishable as a zone of high surf ace density in all star Nearest Neighbor counts. Other patches in the elds that coincided with the Neb ula areas also presented high densities b ut when we applied the IRX color selection, NGC 2237 w as the only one –besides NGC 2244— that w as conrmed to coincide with a cluster A fourth group designated as RLE10 is located North of NGC 2244, and although it has a v ery lo w surf ace density compared to the rest of the clusters, it contains 13 IRX sources, from which at least six ha v e colors suggesti v e of v ery lar ge infrared e xcess, while the re gion has a small e xtinction v alue. 4.2.7 Pr operties of Clusters W e analyzed each cluster indi vidually isolating appropriate sub-re gions that v aried roughly from 25 to 120 arcmin 2 depending on the apparent e xtension of the clusters in the maps. This w ay we were able to determine cluster structures with total areas of 6 to 60 arcmin 2 F or embedded clusters in the Molecular Cloud re gions we calculated the e xtension of a cluster as the area A P inside the polygon dened by the 0.2 ( 0 ) 2 contour in each analysis box, and consider as potential members all of the sources (IRX and non IRX) do wn to K=17.25 inside it. Equi v alent radii, R eq can be dened as p A P = p and can be considered as standard estimates of the total e xtensions of clusters. In T able 4–1 we present for each cluster its center coordinates (as dened from equation 2), core radii (equation 3) and equi v alent radii. W e also sho w the number of IRX sources to K < 15 : 75 and the corresponding fraction it represents. The IRX percentages in the neb ula clusters NGC 2244 and NGC 2237 are intrinsically smaller than those in the embedded clusters, roughly 10 vs 18-76 percent. This

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64 T able 4–1: Y oung Clusters Rosette Comple x Cluster RA DEC R cor e R equiv N I RX p N I RX a I RX F b ID center J2000 [pc] K < 15 : 75 PL01 97.96 4.32 0.37 1.16 29 5 0.28 PL02 98.31 4.59 0.94 1.46 32 6 0.33 PL03 98.38 4.00 0.32 1.69 80 9 0.44 PL04 98.53 4.42 1.10 1.85 89 9 0.24 PL05 98.63 4.32 0.86 1.31 57 8 0.18 PL06 98.66 4.21 0.73 0.75 13 4 0.52 PL07 98.88 3.98 0.38 0.88 22 5 0.61 RLE08 98.56 4.32 0.99 1.30 49 7 0.33 RLE09 98.78 3.69 0.74 1.49 65 8 0.76 RLE10 97.78 5.27 1.19 1.15 15 4 0.32 NGC 2237 97.59 4.93 1.94 1.91 36 6 0.15 NGC 2244 97.95 4.94 1.56 2.30 62 8 0.12 a Number of IRX stars with 10th Nearest Neighbor densities abo v e 0.2 ( 0 ) 2 b IRX fraction with respect to total number of stars with K < 15 : 75 inside 0.2 ( 0 ) 2 contour is due to an e xpected lo wer rate of disk survi v al in the presence of UV radiation from numerous OB stars ( Dolan & Mathieu 1999 ), as well as disk e v olution in older stars which results in reduced circumstellar e xcess emission. F or the clusters embedded in the Molecular Cloud areas (PL01-PL07, RLE08 and RLE09) the lar ge IRX fractions are suggesti v e of ages of 1 to 1.5 Myr or younger ( Haisch et al. 2001 ; Hillenbrand 2006 ). The core radii, R cor e of the Rosette clusters (see equation 4.3 ) ha v e a range of 0.3 to 2.0, with an a v erage of 0.93 0.48 pc. The equi v alent radii, R eq range from 0.75 to 2.30, with an a v erage of 1.44 0.44 pc. The distrib utions of these size estimates are sho wn in Figure 4–7 W e also sho w in the gure the distrib ution of the R cor e = R eq ratios, which peak at 0.65 0.27 and ha v e in tw o cases (clusters PL01 and PL03) v alues belo w 0.5. The clusters NGC 2244 and NGC 2237 are e xtended and their core radii are too close in v alue to their equi v alent radii, so that we considered them equal. The distrib ution of core radii and core to total ratios is consistent with the study of Ferreira et al. ( 2005 ), and suggests that clusters, in most cases, ha v e a tight center b ut with well e xtended edges.

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65 Figure 4–7: From top to bottom: distrib ution of core radii, equi v alent radii and core to equi v alent radii ratios for the Rosette clusters W e constructed color -magnitude and color -color diagrams, which we sho w in Figures 4–8 to 4–19 In the K vs. H K color -magnitude diagrams we sho w the photometry for all of the stars inside the corresponding 0.2 ( 0 ) 2 contour and mark separately those with infrared e xcess. W e include the ZAMS locus and a PMS e v olution isochrone of 1 Myr as well as e xtinction v ectors corresponding to 3 times the mean v alue h A V i in the cluster analysis box. Stars f alling to the right of the isochrone are af fected by e xtinction to w ards the line of sight of the cluster re v ealing their embedded nature. In the J H

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66 vs. H K color -color diagrams, the same stars are located abo v e the dw arf and giant sequences along the reddening bands, with the IRX sources located to the right of the MS reddening strip. Those objects located at or near the zero age sequences, which in the color -magnitude diagram locate preferentially to the left of the isochrone, are most probably fore ground stars or e v olv ed cloud members that coincide with the line of sight of the clusters. a) b) c) d) Figure 4–8: a) K band image, b) control magnitude diagram, c) color color diagram and d) Radial Density Prole for the area corresponding to cluster PL01. See te xt for e xplanation.

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67 a) b) c) d) Figure 4–9: Same as Figure 4–8 for cluster PL02. The fourth plot in each panel are radial density distrib utions calculated with a method of equi v alent areas (see e.g. Muench et al. ( 2003 )) for all stars do wn to K = 17 : 25 inside the analysis box es and calculated from the cluster centers. In this plots we indicate the core and equi v alent radii calculated from the IRX Nearest Neighbor distrib utions. W ith the e xception of PL02, PL06 and RLE10, which are the clusters with the lo west surf ace densities, the rest present well dened radial proles, which unfortunately due to poor statistics cannot be t successfully to standard King or Plummer cluster models, b ut sho w well e xtended tails that in some cases (e.g clusters PL01, PL04, PL07,

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68 a) b) c) d) Figure 4–10: Same as Figure 4–8 for cluster PL03. RLE09) present well dened secondary b umps suggesti v e of structure. In the case of the neb ula clusters NGC 2244 and NGC 2237 the radial distrib ution proles sho w a slo w decline that implies a ne gligible core peak, and might be suggesti v e of an e xtended structure. Ho we v er the counts in each of the equi v alent areas used to construct these proles are not corrected by background, and as the e xtinction is lo wer in the neb ula, these proles might be sho wing the ef fect of eld contamination.

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69 a) b) c) d) Figure 4–11: Same as Figure 4–8 for cluster PL04. 4.3 The Fraction of Stars in Clusters Under the assumption that the IRX are tracing the correct distrib ution of embedded populations in the Rosette comple x, we can use them to estimate the fraction of stars that belong to clusters. W e made our calculations inside the molecular cloud areas rst, to account for deeply embedded clusters only and then for the whole surv e y area which includes the emer ged clusters located in the Neb ula area. The total number of IRX stars detected in the surv e y is 1169 34, out of which 630 25 stars ha v e NN densities lar ger than the mean, 0.2 ( 0 ) 2 A total of 436 21 stars

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70 a) b) c) d) Figure 4–12: Same as Figure 4–8 for cluster PL05. are contained within the estimated areas of the 9 embedded clusters PL01 to RLE09, which occup y a total of 242 sq. arcmin. The remaining 539 23 stars ha v e local surf ace densities lo wer than the mean, and thus cannot be distinguished from the background eld. The area of the molecular cloud co v ered by our surv e y w as calculated as the one contained inside inte grated intensity contour le v els higher than 15 K km s 1 This area is equal to 2747 sq. arcmin, which means the clusters occup y roughly 9% of the cloud. Inside the molecular cloud areas we counted 124 11 stars with densities lo wer

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71 a) b) c) d) Figure 4–13: Same as Figure 4–8 for cluster PL06. than the mean, and 43 7 are stars with densities lar ger than the mean b ut not associated with the cluster areas. F or background correction purposes, we use a scale f actor equal to the ratio of the non-cluster areas of the molecular cloud to the area of the control elds. Using this f actor we e xpect to see a total of 94 10 eld IRX stars, which lea v es a total of 73 8 IRX sources in the cloud areas that are not associated with a cluster From this, we estimate that the fraction of stars in clusters in the Rosette Molecular Cloud is 86 5%.

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72 a) b) c) d) Figure 4–14: Same as Figure 4–8 for cluster PL07. If we repeat these estimates for the whole area of the FLAMINGOS RMC surv e y 7308 sq. arcmin, we nd that there is a total of 549 23 sources associated with clusters after including NGC 2244, NGC 2237 and RLE10 in the counts. The clusters occup y a total area of 390 sq. arcmin, or 5.3% of the total surv e y areas. The number of IRX stars with densities lo wer than the mean is 539 23, and there are 81 9 IRX stars with high densities b ut no association with a cluster The scaled number of e xpected IRX eld stars from the control elds is 261 17 in this case, which results in a total of 359 19

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73 a) b) c) d) Figure 4–15: Same as Figure 4–8 for cluster RLE08. stars non-associated with clusters and a total fraction of 60 5% of stars associated with clusters in the whole surv e y Another interesting result is that in the case of the embedded cluster population, 208 15 sources are contained in clusters PL04, PL05, PL06 and RLE08, at the “Central Core” of the cloud, which corresponds to 48 3% of the total number of embedded sources. This means that approximately half of the recent births in the RMC occurred at the most dense re gion of the cloud, which coincides with the main zone of interaction with the Neb ula ( He yer et al. 2005 ). If the whole surv e y is considered, then the Central

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74 a) b) c) d) Figure 4–16: Same as g 4–8 for group RLE09. Core clusters plus the clusters in the Rosette Neb ula, NGC 2244, NGC 2237 and RLE10, account for 56 3% of the recent stellar formation in the Rosette Comple x, which suggests that the formation occurred in tw o main episodes which resulted in the generation of the biggest clusters, and then, a number of secondary episodes resulted in the smaller remaining clusters which are distrib uted in the remaining areas of the Comple x.

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75 a) b) c) d) Figure 4–17: Same as Figure 4–8 for cluster RLE10. 4.3.1 Distrib ution of Sour ces with Respect to the Rosette Neb ula It is important to mention that with the possible e xception of RLE10, NGC2237 and NGC2244, all of the clusters are associated with a massi v e molecular clump, which conrms their deeply embedded stage. From this, it is clear that the Neb ula and the Molecular Cloud areas e xpose dif ferent episodes of formation. Also, high density IRX stars in the RMC area are mostly conned to the limits of the cloud, while in the Neb ula area, the stars that trace cluster populations are already e xposed out from the molecular

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76 a) b) c) d) Figure 4–18: Same as Figure 4–8 for cluster NGC 2237. gas, possibly sho wing a more e v olv ed population which e v acuated most of the molecular material in the northern half of the comple x. W e calculated the distrib ution of IRX sources with densities higher than the mean as a function of the distance to the center of NGC 2244. T o do this, we counted the number of IRX sources inside the central parsec of NGC 2244 (11 sources) and then we counted those outside this area in concentric annular wedge sectors with a constant width of 1.0 pc, b ut assuring that these sectors were al w ays contained within the surv e y map areas. W e scaled and normalized the star counts in each wedge area with respect to the area of the

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77 a) b) c) d) Figure 4–19: Same as Figure 4–8 for cluster NGC 2244. rst parsec circle. The result is sho wn in Figure 4–20 and we mark ed in the gure the approximate locations of clusters and main cores of the molecular cloud. In the top panel of the gure, we see ho w the prominence of the Rosette Neb ula clusters indicate the y are the lar gest stellar groups in the comple x. The Molecular Cloud “Rigde”, where clusters PL01 and PL02 are located, and which is the part of the molecular cloud that is in direct contact with the ionization front from the Neb ula, appears to be moderate in its star forming ef cenc y The “Core” or central part of the cloud, which contains most of its mass and which has been suggested as the main re gion

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78 Figure 4–20: T op panel: Distrib ution of IRX stars with NN densities higher than 0.2 ( 0 ) 2 as a function of distance from the center of the Rosette Neb ula (NGC 2244). The counts are made in sectors of 1.0 pc in length and counts in each sector ha v e been scaled and normalized to the area and counts in the central 1.0 pc circle in NGC 2244. Labels indicate the approximate locations of clusters described in this paper as well as the main re gions' of the comple x. Bottom panel: equi v alent distrib ution only for sources non associated with cluster areas. of interaction between the molecular and atomic hydrogen clouds (see Celnik 1985 ; Cox et al. 1990 ), and where the clusters PL04, PL05, RLE08 and PL06 are located, seems to be carrying most of the cluster mode production, enhanced by the presence of

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79 cluster PL03, which is ho we v er located in a separated sub-cloud b ut at the same radial distance. Clusters PL03, PL04, PL05, PL06 and RLE08 account fot 58% of the total cluster population. At the “Back Core” of the cloud, there are tw o clusters, PL07 and RLE09 which, although smaller than those in the central core, still ha v e signicant e xtensions. P articularly RLE09 has an e xtraordinary number of young sources despite its location well be yond the interaction front of the Neb ula. F or these clusters, it is possible that a mechanism dif ferent than triggering by interaction with the e xpanding HII re gion need to be proposed. In the bottom panel of Figure 4–20 we repeated the counting b ut only considering stars in the wedges that are located outside of the clusters. The sources not associated with clusters were dened as those located at least tw o cluster radii a w ay from each cluster center The scaled and normalized counts for these stars are of course much smaller b ut it can be seen that there are tw o major zones of along the wedge where noncluster young sources accumulate: the rst one is the area between the NGC 2244 and the “Cloud Ridge”, and the second one is the re gion in between the cloud “Central” and ”Back” cores. 4.3.2 A Case f or a Distrib uted P opulation? Noticing ho w there is a signicant number of sources not associated with clusters in the re gion between the central and back cores of the cloud, we used Field 09 of the surv e y for a separate analysis. This eld lies precisely to the south of the cloud “Central Core” and north of the cloud “Back Core”. The seeing and observing conditions for this eld were particularly good, and 93% of the original area of the eld w as k ept after the polynomial correction. The a v erage scatter of colors do wn to K = 17 : 25 remains belo w an acceptable 0.109 mag across the whole eld, probably because the southeastern quadrant, which for other elds presents high stellar prole distortions, o v erlaps in this case with the good quality northwestern quadrant of the Gap 1 eld. The weighted a v eraging

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80 Figure 4–21: Cumulati v e Counts of IRX sources in eld 09 of the surv e y The solid white bar histogram indicates the number of IRX sources in the eld do wn to brightness limits of K=15.75,16.25,16.75 and 17.25 mag inside the area of best PSF photometry quality delimited by a circle centered on pix el (3000,2400). The dotted line bar histogram indicates the a v erage of IRX counts inside the same circular area for elds 03, 13, 14 and 15, all located in areas mostly de v oid of strong molecular hydrogen emission. The shaded histogram indicates the scaled number of “e xpected” background IRX sources, a v eraged from counts in the control elds (also in the best PSF circular areas), and adding an uniform e xtinction v alue of 5.0 mag. of photometry during the mer ging in this o v erlap area helped to reduce the total color scatter No clusters were found in this eld, b ut a group of 5 IRX sources with NN densities higher than the mean coincide well with a molecular cloud lament located southwest of the PL06 cluster re gion, as seen in Figure 4–5 The absence of clusters or high density re gions complicates the separation of background IRX sources with the NN method, so for this particular e xperiment we simply counted the number of IRX sources in the eld do wn to brightness limits of 15.75, 16.25, 16.75 and 17.25 mag. In order to

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81 Figure 4–22: Examples of lo w density (distrib uted) formation in eld 09 of the sur v e y The blue arro ws indicate the positions of stars with infrared e xcess. Some of the stars with e v en more red colors do not ha v e reliable J or H band photometry ha v e mak e a conserv ati v e estimate, ho we v er we limited the counts to a circular area centered on pix el position (3000,2400), which delimits the area of best PSF photometry quality after the polynomial correction. W e compared these counts with the e xpected number of eld IRX sources at these same limits in an equi v alent area of the control elds, adding a mean e xtinction le v el of 5.0 mag to account for stars that ha v e IRX colors after reddening. A second control counting w as done by a v eraging the number of observ ed IRX sources in the equi v alent circular areas of elds 03, 13, 14 and 15, all located in re gions of the surv e y where the emission from the moleclular cloud w as lo w The results are sho wn in gure 4–21 in the form of cumulati v e count histograms. From the gure, is clear that eld 09 surpasses signicantly the e xpected number of IRX from the background: the counts are 4.8, 5.6, 3.5 and 2.2 times lar ger than the background eld at the respecti v e brightness limits used, and also are higher than the a v erage of f-cloud elds by f actors of 1.3, 1.5, 1.6 and 1.8. This suggests that the re gion of the cloud observ ed in eld 09 has a signicant e xcess in of young sources non associated with clusters. Furthermore, f alse JHK color images (see Figure 4–22 ) of selected zones of the eld sho w a number of highly reddened sources. Some of these sources ha v e

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82 thin neb ulosities, and coincide or are located close to stars with near infrared e xcess emission. The reason why some of the most reddened sources in the images do not present infrared e xcess is because their J band photometry is null or decient and impedes ha ving tw o colors to estimate an e xcess. These pieces of e vidence, along with the distrib ution of non cluster sources in Figure 4–20 suggest that the RMC may be forming stars in a distrib uted or v ery lo w density cluster mode, which w ould be simultaneous with the cluster formation. Unfortunately this hypothesis cannot be conrmed yet because the quality of the photometry in our surv e y has lar ge uncertainties that may put this result into question, especially at the f aint end. Spectroscopic observ ations along with mid infrared imaging are necesary to conrm the numbers of IRX sources in the elds considered, and in case the counterparts are positi v e, then it w ould be necesary to in v estigate a mechanism for their formation which w ould need to e xplain and independent and possibly simultaneous formation of lo w density groups of stars in the cloud laments. 4.4 Discussion and Futur e W ork The cluster population in the Rosette Comple x accounts for 60 to 87% of the present day stellar production, depending on the inclusion of neb ula clusters areas. F or the Molecular Cloud re gion, the 87% fraction w ould be in f air agreement with the results from Lada et al. ( 1991b ) and Carpenter ( 2000 ) for near Giant Molecular Clouds, in which it w as also found that a major fraction (50 to 96%) of the young stellar population w as produced in clusters. In the case of the Rosette, the central core of the cloud, which is also its most dense and massi v e re gion, contains four clusters: PL04, PL05, PL06 and RLE08 which accounts for 48% of the present day formation. Ho we v er although the tw o lar gest clusters in the central core of the Rosette cloud, PL04 and PL05, are the richest among the embedded population, other rich clusters are distrib uted along the remaining area of the cloud: if clusters PL03 and RLE09 are combined with PL04, PL05 and RLE08 as the

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83 richest in the whole cloud, the y account for 77% of the total number of embedded stars, lea ving the other 23% to clusters which are signicantly less numerous. This result is similar to what has been found in other clouds, for e xample, Carpenter et al. ( 2000 ) found that in the comple x W3/W4/W5 half of the population is contained in the v e richest clusters which are associated with the most massi v e clumps; also Lada et al. ( 1991b ) found that one cluster NGC 2024, which coincides with the densest core of the cloud L1630, contains approximately 50% of the total number of stars. This suggests that the distrib ution of cluster mass in molecular clouds follo ws closely the distrib ution of molecular gas clumps, with a small number of massi v e clumps containing most of the total clump mass of the cloud, just as it w as found for Orion B ( Lada 1992 ) and the Rosette itself ( W illiams et al. 1995 ). Please refer to Chapter 6 for a separate discussion on ho w the cluster mass distrib ution compares to the clump mass distrib ution. The clusters in the Rosette appear to be, in the a v erage, slightly lar ger than those reported in other clouds: the cluster equi v alent radii we found range from 0 : 88 to 2 : 3 pc with an mean v alue of 1.44 and a median of 1.46, while in the compilation of Lada & Lada ( 2003 ), clusters run from ( 0 : 3 to 3 : 8 pc in radii, b ut ha v e a mean of 0.80 and a median of 0.62. While for most of the Rosette clusters we observ e a compact cluster core as dened from our NN analysis, the clusters are e xtended well be yond it, as sho wn by the distrib ution of core to e xtended size ratios, in which 5 out of 9 embedded clusters ha v e core to total ratios belo w 0.6. The tw o clusters in the Neb ula area are not well dif ferentiated which suggests that the y are already e xpanded as a result of their early dynamical e v olution, and the particular case of NGC 2244 is special because its lar ge e xtension (2.3 pc) is only comparable to a fe w clusters in literature that surpass the 1.5 pc limit, p.ej. MonR2, Gem4, NGC2282, the T rapezium/ONC cluster Lar ge clusters could be suggesti v e of either a v ery rapid e xpansion, or a mechanism of propagation of formation to w ards the edges of clusters enhanced by rapid emer gence

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84 and dispersion of stars from their parental cores. The second idea could ha v e obser v ational support by the presence of apparently younger groups lik e RLE08 and PL06 located at the edge of the well spread PL04-PL05 Central Core area. Are these groups formed by propagation? A similar ef fect w as observ ed by Lada et al. ( 2000 ), in which protostar candidates in the T rapezium cluster appear to be more conned than older stars to w ards the northwestern edge of the cluster The clusters PL04 and PL05 are signicantly more e xtended than other embedded clusters in the RMC, and their stars are brilliant enough to the point that some can be detected in optical DSS images. These clusters could be slightly older than those located in their periphery which w ould indicate that after the formation of clusters in the Neb ula area, molecular material w as collected and collapsed rst in the core of the cloud, leading to the formation of lar ge associations rst, follo wed by a more modest episodes which lead to the birth of the peripheral clusters. The Central Core of the cloud is also the re gion where the interaction of shock front from the e xpanding HII re gion and the molecular cloud is onset, as observ ed by other authors ( Cox et al. 1990 ; W illiams et al. 1995 ). The interaction of the Rosette Neb ula and the molecular cloud areas, could indeed trigger the rapid, non stopping e v olution of the cloud, which w ould be in agreement with the model of sequential formation of Elme green & Lada ( 1977 ), at least for the central part of the RMC. As the cloud is already f ar from equipartition as sho wn in WBS95, then the cloud could actually still forming stars or e v en condensing in some areas that will host ne w episodes of formation, for e xample the clusters located in the “Back Core” of the Cloud. The hypothesis of induced formation is not to be discarded yet, b ut to conrm it will require e xtensi v e modeling of the dynamical properties of the cloud and its interaction with the e xpanding neb ula, as well as a careful spectroscop y study to determine the age dif ferences of stars along the cloud. Another interesting result from our analysis is the e xistence of an apparently signicant number of IRX stars across a lamentary zone of the cloud which is de v oid of

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85 cluster formation. The IRX counts in this area surpass the e xpected counts in the control elds and in elds located at zones of lo w molecular gas emission, and could represent direct e vidence of a distrib uted mode of formation. If all of the IRX stars located in the molecular cloud b ut outside of the cluster boundaries were members of the comple x, then the y w ould account for approximately 18.5% percent of the young population (81 out of 436 sources). The nature of such a population is unclear as discussed by Carpenter ( 2000 ), because it w ould permeate through the whole molecular cloud, around and in between re gions of cluster formation. In the w orse case scenario these objects could not be members of the comple x, and could be located in the fore ground or background of the cloud. Of course, the e xistence of a small fraction of distrib uted stars w ould not change the main result of this study which is that the Rosette comple x is dominated by cluster formation. Ho we v er in the case that at least a small fraction of these lo w density stars belong to the Comple x and ha v e ages similar to those of the embedded clusters, it w ould be interesting to consider the possible mechanisms that could allo w simultaneous cluster and loose star formation. F or instance, we kno w that the physical mechanism of cloud fragmentation leading to clusters and isolated stars should not be intrinsically dif fer ent. Ev en in clouds where an loose formation mode seems to be preferred lik e T aurus, there is e vidence for high neighbor densities around the lar gest zones of formation (see G omez et al. 1993 ), and the observ ed sizes of star forming cores in the area are similar in size to the neighbor separations, suggesting that the formation of indi vidual cores that form stars one at a time w as not the most common case. Ho we v er the mechanism of fragmentation of a massi v e cluster forming core that leads to the formation of a group of stars with a mass spectrum that matches the observ ed stellar IMF remains a mystery In this sense, if a lo wer scale formation mode is included, and added up to the cluster populations, it w ould be easier to think of a mode of formation in which stars of all masses are

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86 formed in v arious modes, in clusters of dif ferent sizes and then mer ged into lar ge clusters which w ould ha v e the observ ed mass spectra, as suggested by Bonnell et al. ( 2001 ).

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CHAPTER 5 OBSER V A TIONS OF CLUSTER DENSE GAS ENVELOPES 5.1 Intr oduction In this chapter we present the results of our IRAM observ ations of the molecular gas en v elopes of eight of the embedded clusters embedded in the RMC. The goal of this complementary project is to map the emission of dense gas around the clusters in order to study the structure, chemistry and dynamics of their parental core remnants and to ha v e a better understanding of the embedded nature of these clusters, particularly of their relation to the local en vironment. The relation between embedded star clusters and their local en vironments is highly rele v ant in the global picture of star formation: the embedded phase of stars is merely transient, with timescales shorter than those of dynamical e v aporation. This is due to the acceleration of the core gas e xpulsion by means of the stellar winds themselv es. As a result of the rapid dispersion of parental gas, which at the early phases of e v olution of a cluster is the primordial ”glue” that k eeps stars from becoming gra vitationally unbound ( Lada et al. 1984 ), only 10% or less of embedded clusters in our galaxy survi v e as bound structures after emer gence from their parental cores. Emer gence and dispersion occurs typically within the rst 10 Myr of life of clusters and as a result of this, stellar clusters ha v e a high inf ant mortality rate (see Lada & Lada 2003 ). The relation between a young cluster and the remnant molecular gas remnants is also crucial to understand mechanisms lik e mass se gre gation ( Bonnell & Da vies 1998 ) and sequential star formation (SSF) ( Elme green & Lada 1977 ). SSF is particularly important, because the winds of one group of young stars could be directly responsible for altering the dynamics of its remnant medium, leading to the formation of the ne xt stellar birth episode ( Langer et al. 1996 ; Sugitani et al. 1995 ). 87

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88 Also, we w ould lik e to in v estigate ho w the properties of indi vidual clusters depend on the conditions of their forming en vironments. Is star formation more dependent on local (clump scale) or global (cloud scale) conditions? The RMC is an e xcellent re gion to in v estigate this question, as it hosts a signicant number of deeply embedded clusters within a relati v ely small area. All of the RMC clusters are associated with a massi v e molecular clump, as sho wn by W illiams et al. ( 1995 ) and Phelps & Lada ( 1997 ). At the local le v el, molecular clumps associated with cluster formation ha v e not been studied at a resolution capable of rendering their structure at scales comparable to the distances between cluster members. F or e xample, the beam size used to construct the 13 CO map of W illiams et al. w as 1 : 0 7, or 0.79 pc –barely the equi v alent radii of some of the clusters as estimated from the near -IR surv e y The cluster forming clumps identied in the mentioned study ha v e not been mapped with molecular tracers sensiti v e to higher density material either In general, CO isotopes are good to trace the lar ge scale structure of the cloud, at densities of 10 2 10 3 cm 2 b ut to observ e the morphology of star forming cores in detail, we require the use of molecular tracers with critical densities of 10 4 10 5 cm 2 T o resolv e the structure of the molecular gas at sub-cluster le v els, with high critical density tracers, will allo w us to understand the physical processes in v olv ed in the emer gence of the stars from their parental cores. At the global le v el, the clusters are located at dif ferent distances from the HII re gion, b ut those that are located where the cloud is interacting with the neb ula ionization front are particularly interesting. The RMC clumps identied by W illiams et al. ( 1995 ) that ga v e rise to the clusters identied by Phelps & Lada ( 1997 ) had lar ger e xcitation temperatures and mean densities, and that their star forming nature w as denitely correlated to the closeness to the HII re gion, b ut other questions remained unanswered, lik e what is the inuence of the HII re gion into other properties of the clumps directly related to their product stars, lik e the star formation ef cenc y (SFE), a crucial parameter for the dynamical survi v al of a cluster ( Lada et al. 1984 ).

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89 T able 5–1. Rele v ant Properties of Rosette Clusters ID l b N ? (emb) a M(emb) b S Imax RX h A V i c (center) [M ] [ ( 0 ) 2 ] [mag] PL01 206.849 -2.375 55 23 17.5 9.3 1.8 PL02 206.770 -1.941 35 15 12.2 7.2 2.9 PL03 207.326 -2.151 139 77 28.8 9.2 2.5 PL04 206.995 -1.811 130 51 23.3 8.5 2.9 PL05 207.157 -1.782 117 48 40.8 9.7 3.6 PL06 207.268 -1.806 23 39 13.1 10.3 3.5 PL07 207.574 -1.717 43 16 15.8 10.5 2.5 RLE08a 207.125 -1.844 43 20 24.5 11.8 3.1 RLE08b 207.125 -1.844 31 15 24.2 10.5 2.9 a Number and total mass of stars enclosed inside the 5.0K km s 1 contour of 13 CO(2-1) do wn to K < 15 : 75 mag c T otal stellar mass inside the 5.0K km s 1 contour of 13 CO(2-1) calculated for a 1 Myr PMS model and the IMF of the T rapezium Cluster from Muench et al. ( 2000 ) c Median Extinction from Near -Infrared Colors in 13 CO(2-1) observ ed area. 5.2 Obser v ations and Data Reduction Observ ations were done with the 30m telescope at the IRAM f acilities atop Pico V eleta in Granada, Spain during July 25-29th under e xcellent weather and instrument conditions. Our tar gets were the areas corresponding to the clusters PL01 to PL07 and RLE08A. W e obtained high resolution millimeter emission maps in a suite of molecular lines: three at the 1.3mm band, CO(2-1), 13 CO(2-1) and C 18 O(2-1) (with rest frequencies of 230.5, 220.5 and 219.5 GHz respecti v ely) and three at the 3mm band, CS(2-1) HCO + (1-0), and N 2 H + (1-0) (with rest frequencies of 97.9, 89.2 and 93.2 GHz respecti v ely). The 1.3mm molecules aim to trace gas temperature (CO), as well as density of thin gas and gas associated with dust ( 13 CO and C 18 O), while the 3mm tracers were observ ed to trace dense star forming gas with dif ferent chemical pathw ays. These observ ations were based on an e xploratory surv e y of the cluster areas carried out with the FCRA O 14m telescope in March 2001. These observ ations had

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90 a beam size of 50 00 (about 0.4 pc) and sho w the e xistence of signicant, high critical density gas emission. At FCRA O the cluster areas were mapped in four tracers: CS(2-1) HCO + (1-0), 13 CO(1-0) and C 18 O(1-0), b ut the resolution w as not enough to map in detail the structure of the clumps within the cluster scales. The FCRA O maps, ho we v er w ork ed ef fecti v ely as nding charts for the re gions of emission at the cluster core areas, and combined with our FLAMINGOS deep near infrared maps were k e y to determine the re gions of dense gas emission and thus permitted us to use v aluable IRAM observing time ef ciently The IRAM telescope has a main instrument equipped with 8 recei v ers. 4 recei v ers at a time can be tuned simultaneously to indi vidual frequencies. W e used the A100-B100A230-B230 conguration which allo wed us to observ e simultaneously in the 97-115.5 and 197-266 GHz ranges. The typical system temperatures for our setup were 120 and 180 K for the A100 and B100 recei v ers, and 500 and 850 K for the A230 and B230 recei v ers respecti v ely The beam sizes of the detector at 1.3 and 3mm are 12 and 25 : 00 respecti v ely although for the analysis the maps were con v olv ed to 15 and 30 : 00 resolutions to impro v e the signal to noise ratios. W e used the On-the-Fly (O TF) observing method, whereby the antenna scans continuously across a specied re gion and the outputs from the recei v ers are read out on timescales shorter than half the beam transit time, producing ne sampled maps. The O TF scanning of the cluster areas w as done in ro ws separated by 6 00 (1/2 beam at 1.3mm), with a scan rate of 2 00 /sec and readouts e v ery 4 00 /sec, resulting in sampling at 1/3 of the beam size. W e did an of f (reference) position subscan e v ery tw o ro ws for a total reference time of t on p ma psize = 2 F or purposes of calibration, re-focusing and pointing checks, we mo v ed the antenna to w ards Saturn e v ery 3 hours of ef fecti v e observing time. Also, before each O TF scan, we performed a 'cold' (calibration) subscan to dene zero le v els.

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91 Figure 5–1: Maps of IRX sources surf ace density and observ ed molecular emission at the locations of se v en Rosette Clusters and a cluster -less core. The emission of CO(2-1) is contoured as T peak while the other 5 molecules are contoured as inte grated intensity The contour base and step for the contour le v els in each tracer are indicated at the top of each map, in the format base:step F or the IRX density maps the contour le v els indicate surf ace densities of 0.2, 0.6, 1.0, 1.5, 2.0, 2.5, 5.0, 7.5, 10.0, 15.0, 20.0 and 30.0 ( 0 ) 2

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92 Figure 5–2: Same as gure 5–1 b ut for cluster PL02

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93 Figure 5–3: Same as gure 5–1 b ut for cluster PL03

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94 Figure 5–4: Same as gure 5–1 b ut for cluster PL04

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95 Figure 5–5: Same as gure 5–1 b ut for cluster PL05

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96 Figure 5–6: Same as gure 5–1 b ut for cluster PL06

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97 Figure 5–7: Same as gure 5–1 b ut for cluster PL07

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98 Figure 5–8: Same as gure 5–1 b ut for cluster RLE08A. Notice that the dense gas emission for this cluster dene tw o separated contours, which we labelled as RLE08A1 and RLE08A2.

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99 PL01 PL02 PL03 PL04 Figure 5–9: Near -Infrared Extinction maps to w ards the areas mapped in 13 CO(2-1)for RMC clusters PL01 to PL04. The infrared maps ha v e been con v olv ed to a resolution of 30 00 same as the IRAM resampled maps. 13 CO(2-1) emission is indicated by solid contours, in le v els of 5 K km s 1

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100 PL05 PL06 PL07 RLE08 Figure 5–10: Near -Infrared Extinction maps to w ards the areas mapped in 13 CO(2-1)for RMC clusters PL05 to RLE08. The infrared maps ha v e been con v olv ed to a resolution of 30 00 same as the IRAM resampled maps. 13 CO(2-1)emission is indicated by solid contours, in le v els of 5 K km s 1

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101 W e originally aimed to obtain maps with co v erage of at least 4 : 0 4 : 0 for e v ery cluster re gion, b ut depending on the strength of the signal obtained to w ards each cluster in e xploratory pointings, we increased some maps to a lar ger area co v erage by making tw o partially o v erlapping maps, or reduced the scan area of single maps as necessary (please refer to T able B–1 Appendix B for co v erage details). The signal for the weak er emission molecules, N 2 H + (1-0) and C 18 O(2-1), w as lo w to the point that the O TF scans could only be constructed around the re gions of strongest signal detection, resulting in small maps that might not co v er the entire area of the clump, or at least the same area as the other tracers. In tw o cases, PL02 and RLE08B, the signal to noise w as so weak that N 2 H + (1-0) could not be detected. In the case of PL02, neither C 18 O(2-1) nor N 2 H + (1-0) ga v e a good signal and the resultant, v ery noisy map w as discarded. W e selected to observ e the tracers with stronger signals rst, tuning the detector with HCO + (1-0) and CS(2-1) in the A100 and B100, and 13 CO(2-1) and CO(2-1) in the A230 and B230 recei v ers respecti v ely Once the eight clusters were observ ed in these tracers, we proceeded to observ e elds with N 2 H + (1-0) in the A100 and B100 and C 18 O(2-1) in the A230 and B230 recei v ers. In a couple of cases, the second, weak er tracer conguration scan w as repeated to increase the quality of the maps in those clusters that were observ ed the rst time under belo w a v erage conditions. Reduction of the data w as done by obtaining rst the system temperatures (zero points) from a cold scan, then scaling the spectra to these, and then subtracting the a v erage of 2 adjacent reference subscans from a O TF scan to obtain a 'pure' signal. Once the scans were calibrated, we remo v ed baselines adjusted mostly to rst order and then indi vidual spectra were processed with standard CLASS routines to create data cubes. Finally standard GILD AS softw are routines were used to put the maps into FITS le format.

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102 T able 5–2. Clump properties for Rosette Clusters line l o f f b o f f R cl um p h v i D v R T d v M vir N N( H 2 ) M L T E (arcsec) (pc) (km s 1 ) (K km s 1 ) (M ) (cm 2 ) (M ) CLUSTER PL01. T e x = 17.3 K 13 CO(2-1) 3.42 -2.92 0.73 13.8 2.07 1.47e+04 654 1.43e+16 7.61e+21 284 HCO + (1-0) 1.97 -0.263 0.72 14.2 1.41 2.46e+02 299 5.10e+12 7.11e+21 256 CS(2-1) 28.6 19.5 0.60 13.7 1.43 1.94e+02 257 8.80e+12 4.95e+21 125 N 2 H + (1-0) 22.5 26.4 0.50 12.7 1.10 0.23e+02 128 3.85e+11 3.39e+21 59 C 18 O(2-1) 2.60 42.6 0.60 13.7 1.74 1.68e+03 376 1.60e+15 5.45e+21 135 CLUSTER PL02. T e x = 20.9 K 13 CO(2-1) 24.0 10.6 0.76 13.8 1.82 1.26e+04 527 1.32e+16 6.40e+21 259 HCO + (1-0) 34.9 7.77 0.76 15.1 1.42 2.12e+02 323 5.05e+12 6.12e+21 247 CS(2-1) 39.7 -1.73 0.61 13.8 2.66 1.14e+02 911 5.79e+12 3.79e+21 99 CLUSTER PL03. T e x = 27.2 K 13 CO(2-1) -25.8 -8.21 0.81 15.3 1.42 1.44e+04 338 1.73e+16 1.12e+22 410 HCO + (1-0) -5.40 -6.95 0.83 15.7 1.71 2.38e+02 508 6.97e+12 1.40e+22 465 CS(2-1) -79.7 -4.88 0.64 15.2 0.874 1.44e+02 102 8.83e+12 9.67e+21 151 N 2 H + (1-0) -80.6 -2.87 0.44 14.1 0.956 0.14e+02 83 3.35e+11 1.15e+21 37 C 18 O(2-1) -76.9 -16.8 0.60 14.9 1.38 1.58e+03 238 1.67e+15 1.61e+21 135 CLUSTER PL04. T e x = 36.4 K 13 CO(2-1) -38.6 -9.94 0.93 15.2 1.79 3.31e+04 625 3.25e+16 1.12e+22 679 HCO + (1-0) -79.5 88.0 1.1 15.4 1.70 5.40e+02 640 2.03e+13 1.40e+22 1090 CS(2-1) -74.0 71.6 0.85 15.2 1.39 2.90e+02 344 2.22e+13 9.67e+21 488 N 2 H + (1-0) -82.2 93.6 0.38 14.3 1.37 0.13e+02 148 3.79e+11 1.15e+21 11 C 18 O(2-1) -81.5 80.0 0.45 15.0 1.30 9.21e+02 158 1.11e+15 1.61e+21 23 CLUSTER PL05. T e x = 12.9 K 13 CO(2-1) 89.9 -30.18 0.75 11.6 1.79 8.60e+03 499 8.14e+15 8.22e+21 318 HCO + (1-0) 77.6 -20.64 0.55 12.1 1.21 0.64e+02 169 1.08e+12 4.45e+21 95 CS(2-1) 79.7 -33.60 0.66 11.8 0.92 1.51e+02 118 5.93e+12 6.30e+21 193 N 2 H + (1-0) 70.0 -14.46 0.43 11.3 0.97 1.22e+02 84 1.65e+11 3.04e+21 39 C 18 O(2-1) 96.6 -35.64 0.43 11.7 1.15 5.36e+02 118 5.29e+14 3.68e+21 46 CLUSTER PL06. T e x = 16.4 K 13 CO(2-1) -13.6 -8.21 0.90 12.5 2.39 2.11e+04 1080 2.00e+16 1.28e+22 726 HCO + (1-0) -9.66 -13.6 0.86 13.4 1.52 2.86e+02 415 5.70e+12 1.21e+22 627 CS(2-1) -19.9 -21.0 0.84 12.6 1.47 3.95e+02 380 1.74e+13 1.14e+22 554 N 2 H + (1-0) -17.3 -14.5 0.68 11.9 1.56 1.12e+02 346 1.76e+12 7.58e+21 243 C 18 O(2-1) -23.4 -18.9 0.78 12.5 1.50 3.38e+03 365 3.22e+15 9.43e+21 399 CLUSTER PL07. T e x = 17.8 K 13 CO(2-1) 7.57 2.91 0.90 12.2 1.98 2.11e+04 734 2.06e+16 1.29e+22 722 HCO + (1-0) 37.4 25.1 0.62 11.1 0.87 2.86e+02 98 1.96e+12 6.44e+21 172 CS(2-1) 10.1 14.9 0.90 12.2 1.53 3.95e+02 439 2.59e+13 1.24e+22 694 N 2 H + (1-0) 36.8 31.5 0.60 11.6 1.84 1.12e+02 425 2.07e+12 5.80e+21 144 C 18 O(2-1) 8.47 -17.6 0.65 12.0 1.57 3.38e+03 336 1.95e+15 7.28e+21 216 CLUSTER RLE08a T e x = 13.4 K 13 CO(2-1) 27.0 -4.71 0.81 10.2 1.49 1.40e+04 377 1.32e+16 1.18e+22 544 HCO + (1-0) 5.25 -1.74 0.65 9.94 1.08 0.83e+02 159 1.21e+12 6.95e+21 207 CS(2-1) 67.0 66.8 0.83 10.7 1.75 0.37e+02 531 1.34e+13 1.23e+22 584 N 2 H + (1-0) -63.2 -207. 0.52 10.7 0.85 0.13e+02 79 1.52e+11 4.57e+21 87 C 18 O(2-1) -10.7 29.4 0.81 10.6 1.28 3.10e+03 274 3.47e+15 1.14e+22 512 CLUSTER RLE08b T e x = 9.9 K 13 CO(2-1) -27.3 -102. 0.60 15.7 1.17 7.76e+03 174 7.35e+15 5.74e+21 145 HCO + (1-0) -22.0 -111. 0.60 15.8 0.93 1.01e+02 108 1.77e+12 5.76e+21 144 CS(2-1) -31.1 -92.6 0.53 15.8 0.85 1.11e+02 81 4.42e+12 4.52e+21 90 C 18 O(2-1) -21.6 -90.6 0.54 15.8 0.85 1.04e+03 81 1.02e+15 4.51e+21 91

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103 The nal resolutions of the maps are 5.5 00 for the 1.3mm lines and 12.25 00 for the 3mm lines after Nyquist sampling. The typical RMS noise v alues are 0.31 and 0.12 K km s 1 for the 1mm and 3mm lines respecti v ely 5.3 Analysis and Results 5.3.1 Pr esentation of the Data Figures 5–1 to 5–8 sho w for e v ery cluster re gion, the resultant maps in contours of inte grated emission ( R T d v ) in K km s 1 e xcept for CO, which is plotted as peak temperature (T peak ) instead. W e also included a NIR image of the cluster and contours of IRX sources surf ace density The molecular gas emission contours in the gures were chosen to depict clearly the general morphology of the clump, b ut for analysis we selected the clumps as the re gions enclosed in a constant le v el contour (see section 5.3.3 ). The infrared e xcess contours sho w only the infrared e xcess surf ace density for le v els greater or equal to 0.2 ( 0 ) 2 as measured with the nearest neighbor analysis technique described in section 4.2.1 The maps sho w in general a v ery good correlation between the morphology of the emission contours and the distrib ution of stars in the cluster cores. Ho we v er in some cases, the molecular gas en v elopes appear to enclose the cluster cores completely (e.g. CS(2-1) emission in cluster PL06, Figure 5–6 ), characteristic of a deep le v el of embedeness, while in others some gas seems to be swept out from specic re gions of the cluster suggesti v e of partial emer gence from the gaseous clump (e.g. 13 CO(2-1) emission in cluster PL03, Figure 5–3 ). A brief description of the most interesting characteristics of each cluster map suite can be found in Appendix C In table 5–1 we present data rele v ant to this chapter from the NIR photometry analysis of chapter 4 W e included the number of embedded stars N emb dened as the number of stars counted inside the 5.0 K km s 1 contour of 13 CO(2-1) the peak of local infrared e xcess stellar density of the cluster S I RX max and the total embedded stellar mass M emb which w as calculated with the models of Muench et al. ( 2000 ). These

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104 models allo wed us to estimate a luminosity to mass con v ersion f actor for the observ ed N emb corresponding to a cluster with an age of 1 Myr in the PMS e v olution models of D'Antona & Mazzitelli ( 1997 ). The con v ersion w as calculated by assuming the Rosette clusters ha v e an IMF similar to T rapezium cluster at the distance of the Rosette and ha ving a distrib ution of e xtinction similar to the one observ ed in the map areas. 5.3.2 Local Extinction The clusters in the RMC are deeply embedded in the gas en v elopes, mix ed with a signicant amount of dust. A v erage visual e xtinction A V is traditionally estimated from a direct con v ersion of the column density of a CO isotope, say 13 CO(2-1) into visual e xtinction. An e xample of such a con v ersion is the one from Frerking et al. ( 1982 ): N 13 = 2 : 7 10 15 ( A v 1 : 6 ) cm 2 mag 1 [ 0 < A v < 10 ] : (5.1) Unfortunately for range of visual e xinction is rather short for a typical star forming cloud, where it has been sho wn that e xtinction can easily surpass the 30 mag limit. More recently near infrared colors of stars in the background of dark clouds ha v e been used succesfully to produce e xtinction maps which co v er a lar ger range of magnitudes and sho w a v ery good correlation with the CO column density ( Alv es et al. 1999 ; Lada et al. 1994 ; Ber gin et al. 2002 ). W e are particularly interested in calculating the e xtinction to w ards our molecular maps because we will use those v alues to determine chemical ab undances for the dif ferent observ ed tracers. Using the technique we describe in detail in Chapter 6 we used our FLAMINGOS photometry to construct e xtinction maps with resolutions con v olv ed to those of the indi vidual tracer resampled beam sizes (see section 5.2 ). In gures 5–9 and 5–10 which sho w a good correlation with the inte grated intensity emission contours of 13 CO(2-1)

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105 5.3.3 Calculation of Ph ysical P arameters Properties of en v elopes were calculated directly from the maps using an interacti v e tool program that allo wed us to delineate the boundaries of the clumps from contours of inte grated intensity o v erlayed on FLAMINGOS images of the cluster Contours were constructed after carefully selecting an optimal range of v elocities in the a v erage spectrum of the map. W e made sure that the v elocity channels selected from the a v erage spectrum line e xcluded the line wings, which usually arise from gas outo ws. The results of our indi vidual re gion analysis can be consulted in T able 5–2 F or each cluster area, we included the T e x peak v alue from the CO(2-1) maps, the of fset of each tracer emission peak with respect to the center of stellar density the en v elope size (equi v alent radius) and the v alues of central v elocity and v elocity dispersion for the a v erage spectrum of each map. W e also included virial masses (M vir ), inte grated intensities, tracer column densities, H 2 column densities, and gas (L TE)mass (M L T E ). The a v erage clump size is dened as R cl um p = p A cl um p = p with A cl um p being the area enclosed by a constant minimum contour le v el. F or all the 13 CO(2-1) maps we used the contour at 5.0 K km s 1 for the CS(2-1) HCO + (1-0) and C 18 O(2-1) maps we used the one at 0.5 K km s 1 and for N 2 H + (1-0) we used the one at 0.2 K km s 1 W e dened the central v elocity h v i of a clump as the v elocity corresponding to the maximum of the a v erage spectrum line in each map, and the v elocity dispersion, s v is dened as D v = 2 : 355 s v with D v being the FWHM of the a v erage spectrum line. V irial masses where calculated as M vir = 5 R cl um p s 2v = G assuming that a vir = 1 as in a constant density prole and where G=1/232 in units of solar mass, km/s and pc follo wing Bertoldi & McK ee ( 1992 ). Also, from the a v erage spectra, we included the inte grated intensity R T d v o v er each clump area. This allo wed us to calculate column densities for each clump by using the relation between column density and optical depth from Sco ville et al. ( 1986 ):

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106 T able 5–3. Molecular Line P arameters Line a B a [Debye] [GHz] 13 CO(2-1) 0.11 55.101 HCO + (1-0) 3.30 44.594 CS(2-1) 1.96 24.496 N 2 H + (1-0) 3.40 46.587 C 18 O(2-1) 0.11 54.891 a Obtained from the N ASA Catalog of Molecular Spectroscop y (http://spec.jpl.nasa.go v /) N = 3 k 8 p 3 B 2 e hBJ ( J + 1 ) k T e x J + 1 T e x + hB 3 k 1 e hv k T e x Z t n d v (5.2) which, after applying radiati v e transfer theory and assuming optically thin emission (see Dickman 1978b ) translates into: N = 1 : 737 10 15 B GH z 2Debye e 0 : 048 B GH z ( J + 1 )( J + 2 ) = T e x ( J + 1 ) 2 ( T e x + 0 : 016 B GH z ) Z T R d v (5.3) where the column density is gi v en in units of cm 2 and the inte grated intensity in K km s 1 W e obtained the corresponding v alues for the dipole moment and the rotational constant B for each tracer from the N ASA Catalog of Molecular Spectroscop y 1 These v alues can be consulted in table 5–3 5.3.4 T racer Ab undances W e calculated the a v erage ab undances of each molecular tracer N ( X ) = N ( H 2 ) using our near -infrared e xtinction maps, by calculating the ratio A V = N ( X ) to w ards each line of sight, and the standard v alue N ( H 2 ) = A V = 2 : 06 10 21 ( Spitzer 1978 ; Bertoldi & McK ee 1992 ). The distrib ution of the a v erage ab undance ratios obtained this w ay for each cluster map is sho wn in Figure 5–11 1 http://spec.jpl.nasa.go v/

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107 Figure 5–11: Distrib ution of observ ed ab undance ratios by tracer The labels in each histogram bar indicate the corresponding clump. The a v erage ab undance ratios we obtained are: 1. l o g ( N 13 C O = N H 2 ) = 5 : 79 0 : 13 which is in good agreement with the -5.68 used by Bertoldi & McK ee ( 1992 ), b ut these authors pointed out lar ge discrepancies among the ratio v alues calculated for dif ferent clouds. Ho we v er if we compare to the pre-stellar core v alues from the recent study of Jr gensen et al. ( 2004 ), in which l o g ( N C O = N H 2 ) is -4.85, and taking the 12 C O = 13 C O ratio of 7 : 5 2 : 1 for clouds in M33, we obtain a 13 C O ab undance ratio of -5.72, which is in slightly better agreement with our lo west v alues.

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108 2. l o g ( N C S = N H 2 ) = 9 : 03 0 : 32, which is, within the errors, slightly lo wer than the -8.82 calculated for pre-stellar cores from Jr gensen et al. 3. l o g ( N H C O + = N H 2 ) = 9 : 39 0 : 28, which is lo wer than the v alue of -9.09 calculated for prestellar cores by Jr gensen et al. Our ratios occup y a v ery lar ge range (-9.8 to -8.7), which translates into a lar ge uncertainty 4. l o g ( N N 2 H + = N H 2 ) = 9 : 86 0 : 47, which is lo wer than the -8.57 reported by Jr gensen et al. in the case of pre-stellar cores. Ho we v er much smaller ratios of up to -10.69 were found by T af alla & Santiago ( 2004 ) in other starless cores. Other more comparable v alues v ary from -9.88 to -10.69 for Barnard 68 ( Di Francesco et al. 2002 ) and -9.3 for clouds in the TMC-1 and L134N comple x es ( Ohishi et al. 1992 ). Our range of v alues is v ery ample (-10.6 to -9.0), b ut in terms of uncertainty similar to the dif ferences among the studies cited (-10.7 to -9.3). 5. l o g ( N C 18 O = N H 2 ) = 6 : 57 0 : 20, which is higher than the ranges of -7.52 to -6.88 reported by Di Francesco et al. ( 2002 ) for Barnard 68 and the -6.79 reported by Ohishi et al. for TMC-1 and L134N. Ho we v er 7 out of 8 clusters present ab undance v alues lo wer than the median and can be as lo w as -6.92. 5.3.5 Gas Mass Using the ab undance v alues described abo v e, we calculated then the mass of gas in local thermodynamic equilibrium (L TE) for each clump in each tracer by follo wing the reasoning of Bertoldi & McK ee ( 1992 ): M L T E = 2 m H 1 Y N H 2 N X Z N X d A : (5.4) W ith A being the projected surf ace of the clump. W e used a constant helium fraction o v er the cloud of Y = 0 : 28. In the observ ed units (solar mass, parsecs, grams and densities in cm 2 ), the last equation can be simply e xpressed as M L T E = 6 : 95 10 20 N H 2 R 2cl um p : (5.5)

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109 5.3.6 Clump Sizes Figure 5–12 sho ws the distrib ution of the estimated en v elope sizes, R cl um p for 13 CO(2-1) CS(2-1) and HCO + (1-0). W e compare N 2 H + (1-0)and C 18 O(2-1)separately as the maps are smaller and the resultant clump sizes are systematically smaller too. Figure 5–12: Left: Distrib ution of clump sizes, separated by tracer for CS(2-1) HCO + (1-0) and 13 CO(2-1) Right: Distrib ution of clump sizes, separated by tracer for C 18 O(2-1) and N 2 H + (1-0). The a v erage v alues for the clump sizes from 13 CO(2-1) CS(2-1) and HCO + (1-0) are 0.8 0.1, 0.72 0.13 and 0.74 0.16 respecti v ely which is belo w the estimates for star forming clump sizes in W illiams et al. ( 1995 ) (1.3-4.2 pc), b ut our maps are v ery localized at the centers of near infrared density and comprise only a fraction of the total areas of the W illiams et al. clumps. The sizes we estimate for these local en v elopes ho we v er are not atypical if compared to other studies (e.g. Carpenter et al. 2000 ; Simon et al. 2001 ; Ridge et al. 2005 ; T atematsu et al. 1998 all with typical sizes within 0.2-1.8 pc), and the main results of this study are not af fected by the discrepanc y In the case of N 2 H + (1-0) and C 18 O(2-1) which are e v en more localized, the en v elope sizes are, in the rst case, distrib uted around 0.18 pc and in the second case, distrib uted around 0.42 pc, indicating that the emission for these tracers is much more localized.

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110 5.3.7 Line W idths The distrib ution and comparison of line widths, D v for 13 CO(2-1) CS(2-1) and HCO + (1-0) are sho wn in Figure 5–13 and again, we sho w separately those for N 2 H + (1-0) and C 18 O(2-1). The v alues for most tracers are in good agreement with each other although 13 CO(2-1) has, in the a v erage, slightly lar ger widths than the dense tracers possibly due to line broadening from outo ws. In general, our D v v alue for 13 CO(2-1) agree well with those of W illiams et al. ( 1995 ). Line width is a crucial parameter as M vir is directly proportional to D v 2 so that e v en small dif ferences in the line width estimates increase the uncertainty in the virial mass calculations. W e were as careful as possible as to not include line wings in the v elocity channel selection for the calculation of virial masses b ut depending on ho w broad or ho w noisy indi vidual spectra are, v ariation of the order of 10 1 km cm 1 can translate into M vir v ariations of a fe w to a fe w tens of solar masses, or about less than 10% of the total virial mass. Figure 5–13: Left: Distrib ution of v elocity dispersions, separated by tracer for CS(2-1) HCO + (1-0) and 13 CO(2-1) Right: Distrib ution of v elocity dispersions, separated by tracer for C 18 O(2-1) and N 2 H + (1-0). 5.3.8 Clump Masses Figure 5–14 sho ws the distrib utions and comparisons of M vir v alues obtained for 13 CO(2-1) CS(2-1) and HCO + (1-0) and again, separated those for HCO + (1-0) and

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111 C 18 O(2-1). In gure 5–15 we sho w the corresponding distrib utions for the gas mass, M L T E which are calculated directly from the N ( H 2 ) column density using equation 5.5. Figure 5–14: Left: Distrib ution of virial mass, separated by tracer for CS(2-1) HCO + (1-0) and 13 CO(2-1) Right: Distrib ution of virial mass, separated by tracer for C 18 O(2-1) and N 2 H + (1-0). The a v erage en v elope virial mass v alues obtained for 13 CO(2-1), CS(2-1) and HCO + (1-0)are 501, 263 and 251 M respecti v ely while for N 2 H + (1-0) and C 18 O(2-1)are 244 and 185 M The M vir v alue ranges for the N 2 H + (1-0) and C 18 O(2-1) clumps are closer to those from HCO + (1-0) and CS(2-1) b ut these could be in f act lo wer quotes because of the smaller map areas. Our e xploratory pointings at the telescope did not sho w much more emission outside of the re gions mapped, and we could safely say that the area of the en v elopes for N 2 H + (1-0) or C 18 O(2-1) w as not underestimated by more than a f actor of 2. En v elope virial masses from 13 CO(2-1) are systematically higher than those estimated with the other tracers, which might be a result of the slightly lar ger line widths of the 13 CO(2-1)en v elopes. Ho we v er the distrib ution of line widths is similar for all the other tracers, so that the main uncertainty in the virial mass estimates for those might be the underestimation of en v elope size due to the limit mapping co v erage. The v alues of M L T E are more widely distrib uted, b ut in general present smaller v alues than M vir In Figure 5–16 we can see ho w this results in a signicant number of

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112 Figure 5–15: Left: Distrib ution of L TE mass, separated by tracer for CS(2-1) HCO + (1-0) and 13 CO(2-1) Right: Distrib ution of L TE mass, separated by tracer for C 18 O(2-1) and N 2 H + (1-0). cases with virial-to-gas mass ratios a vir = M vir = M L T E lar ger than one. a vir measures the ratio of kinetic to gra vitational ener gy in the gas clumps, and lar ge ratios indicate that there is a signicant number of cluster en v elopes in the Rosette that are probably not bound by gra vity b ut rather by ambient pressure. Uncertainties in sizes and line widths, might, if at all, translate into lar ger virial masses, resulting in a lar ger fraction of unbound en v elopes.F or 13 CO(2-1), which has the best clump area co v erage, the fraction of bound clumps is 33%, b ut again, 13 CO(2-1) virial masses are af fected by lar ger line widths. The ef fect is less dramatic for CS(2-1) and HCO + (1-0), for which 77% and 56% of the en v elopes present bound ratios respecti v ely F or C 18 O(2-1), the fraction of bound clumps is 38%, and in the case of N 2 H + (1-0), only one clump presents a bound ratio. Ho we v er if we look closely at the ratios from 13 C O which is the tracer that gi v es a better estimate of the total mass of gas associated with the clump, we see that se v en out of nine en v elopes ha v e a vir v alues smaller than 1.7 and the other tw o ha v e v alues belo w 3.0. Systematic discrepancies in the virial and L TE mass calculation methods, specically for 13 CO(2-1) ha v e been pointed out before: for e xample McK ee & T an ( 2003 ), sho wed that reasonable v alues of a vir could be up to 1.4, and Lee et al. ( 1994 ) found e v en lar ger

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113 Figure 5–16: Distrib ution of virial to L TE (gas) mass ratios v alues for the dif ferent tracers. Bound clumps are located to the left of the dashed ( M vir = M L T E ) line. a vir ratios, which range from 2.0 to 6.0. One possible e xplanation for such discrepancies is that L TE might not be a safe assumption in the case of partially e v olv ed clusters lik e ours, because acti v e interaction between winds from massi v e stars and local en vironments in the centers of formation is onset. Unbound clumps could be k ept in local equilibrium by means of e xternal pressure from the inter clump gas. The ambient binding pressure P e x t necessary to maintain the equilibrium in a molecular core, can be estimated indirectly As sho wn by K eto & Myers ( 1986 ) and Simon et al. ( 2001 ), P e x t is a function of mass, size and v elocity dispersion, and its equilibrium v alues can be inferred from the relation between the ratio s 2 = R and the column density N ( H 2 ) both proportional to the surf ace density M = p R 2 The ef fect of e xternal pressure onto the virial equilibrium of a gas core is lar ge for small clumps,

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114 Figure 5–17: The ratios s 2 = R as a function of column density N ( H 2 ) The solid lines e xpress constant binding pressure. where a vir is lar ge, and s 2 = R goes as N ( H 2 ) 1 In lar ge clumps the ratio goes directly as N ( H 2 ) and P e x t is ne gligible. In Figure 5–17 we sho w ho w in the case of the Rosette clusters, 5 en v elopes, with bound ratios close or less than unity (PL03, PL04, PL07, RLE08a and RLE08b) agree with a model in which the e xternal pressure is of the order of P e x t = k 10 5 K cm 3 which is the typical v alue of binding pressures indicated by theoretical studies (e,g. Myers 1998 ). F or the rest of the clumps, the binding pressure w ould need to be up to 1 order of magnitude lar ger 5.3.9 Gas Dynamics The motion of the gas inside the gas clumps w as measured by calculating the v ariation of the v elocity and v elocity dispersion across the indi vidual maps. The observ ed line proles are Gaussian-lik e in most of cases, and inf all motions of dense gas can be discarded due to the absence of self-absorbed proles in the surv e y

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115 Figure 5–18: Distrib ution of v elocity gradients by tracer In gure 5–18 we sho w the distrib ution of v elocity gradient v alues d v = d l for the dif ferent tracers. It is v ery clear that 13 CO(2-1) which traces lo w density gas sho ws little or no v ariation within the clump areas, while the dense gas tracers sho w gradients which v ary roughly from 0.03 to 0.1 km s 1 cm 1 This is suggesti v e of gas motions being more coherent near the en v elope cores where star formation occurs, surrounded by rather turb ulent edges.

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116 T able 5–4. Star F ormation Ef ciencies ( 13 CO(2-1) ) Cluster SFE vir SFE L T E (%) PL01 3.4 7.5 PL02 2.8 5.5 PL03 18.6 15.8 PL04 7.5 7.0 PL05 8.8 13.1 PL06 3.5 5.1 PL07 2.1 2.2 RLE08a 5.0 3.5 RLE08b 7.9 9.4 5.4 The Embedded Stellar P opulation 5.4.1 Star F orming Efciencies Star forming ef ciencies (SFE) were calculated for the 13 C O L TE mass v alue as e = M ? = M ? + M L T E where M ? is the embedded mass of cluster M emb which only considers those stars which were located inside the 5.0 K km s 1 contour of 13 CO(2-1), considered for the en v elope area estimate. The v alues of M emb from 5–1 were calculated by applying the models of Muench et al. ( 2000 hereafter MLL00) to create articial clusters containing the same observ ed number of embedded stars, and supposing an Initial Mass Function (IMF) similar to that of the T rapezium cluster at the distance of the RMC (1600 pc). The ef ciencies obtained are tab ulated in T able 5–4 where for comparison, we also included those calculated using the M vir v alues for 13 CO(2-1). A distrib ution of the SFE L T E v alues is sho wn in Figure 5–19 Our SFE v alues (L TE) v ary roughly from 2 to 16 percent, with 7 out of 9 re gions f alling belo w the 10 percent ef cenc y le v el. Clusters PL03 and PL05 ha v e ef cencies of 16 and 13 percent respecti v ely This is in a reasonable agreement with ef ciencies reported in the literature for other embedded clusters (see table 2 of Lada & Lada 2003 ). Another v ery important point to consider is that our SFE v alues are only local estimates of the true SFE of the cluster forming re gions, because some of the clusters, lik e PL04 and PL05 appear to be partially emer ged. Columns 4 and 6 of T able 5–1 indicate that the number of stars considered to belong to the observ ed ”cores” of the cluster are lar ger with a signicant fraction of the stars observ ed ”outside” of the gas en v elopes,

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117 Figure 5–19: Distrib ution of star formation ef ciencies per cluster calculated from the cluster mass and the 13 CO(2-1) L TE clump mass. The width of the bars (4.4 percent) indicate the standard de viation of the distrib ution. and thus it might be possible that a substantial fraction of the initial mass of star forming gas has already e v aporated. This results in an o v erestimation of the mass of stars in the equation and thus the more e v olv ed is a cluster the more uncertain is its observ ed SFE v alue. 5.4.2 The Gas Stars Connection The images and emission morphology in the dense gas maps suggest that the RMC clusters are in the process of emer ging from their parental cores. One interesting aspect to notice is that the density peaks of the clusters do not necessarily coincide with the peaks of gas density In Figure 5–20 we sho w ho w the projected of fsets might be related to the sizes of the clusters themselv es. The correlation coef cients for the CS(2-1) ,N 2 H + (1-0) and C 18 O(2-1) of fsets are abo v e 0.7 with of fset v alues being v ery similar for the 3 tracers. In the case of HCO + (1-0) the coef cient is 0.4, with noticeable

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118 dif ferences of the order of 30 and 80 arcseconds with respect to v alues for the other tracers for PL01 and 03, respecti v ely The positi v e correlation of projected of fset vs. cluster radius is suggesti v e of both partial emer gence and cluster structure: clusters PL03, 04 and 05, which ha v e the lar gest equi v alent radii and of fsets, ha v e stars that are already visible in deep DSS plates, and their gas cores are small and concentrated to w ard v ery localized groups of hea vily reddened stars. Figure 5–20: Correlation between cluster size and projected emission peak of fsets with respect to the cluster stellar density peaks. The horizontal axis scale is sho wn at the same scale as the v ertical axis for clarity

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119 5.5 Chemical Differ entiation in Cluster En v elopes The apparent density peak of fsets suggest that CS(2-1) and N 2 H + (1-0) could be tracing the most recent star forming condensations of gas. In the case of cluster PL03, we found a group of v ery red stars at the center of density of CS(2-1) This peak is dramatically displaced from the density center of the cluster and from the peak of HCO + (1-0)emission (see Figure 5–21 ). Figure 5–21: An o v erlap of HCO + (1-0) (solid contours) and CS(2-1) (dashed contours) emission to w ards Rosette Cluster PL03. Notice the of fset between the emission peaks. Red dots indicate the position of hea vily reddened f aint sources which are coincident with the CS peak. The contours for both tracers run from 0.25 KKm/s in 0.75 KKm/s steps. The background image is a K band image from FLAMINGOS

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120 In the case of HCO + the emission is more intense near the re gions where dusty gas is displaced and re-condensed (apparently into 3 separate clumps). One possibility is that once H 2 O e v aporates from dust grains due to the presence of ultra violet radiation from partially emer ged stars, OH + and radicals could combine with carbon molecules enhancing the ab undance of HCO + in the presence of dust. CS, on the other hand, might be depleted out of the emer ging areas and reconcentrated in the outskirts where ne w stars are being formed. The chemistry of CS is intimately link ed to CO. When stars form, dust grains reach a temperature where CO can be e v aporated (around 25K), increasing the ab undance of CO in the cores. This ef fect rapidly enhances the formation of CS, which has a similar e v aporation time. Ho we v er in contrast to other molecules lik e HCN or H 2 CO, which cannot be sustained after e v aporation, CS remains ab undant as atomic Oxygen and H 2 O freeze onto grains at small core radii (see Lee et al. ( 2004 )). This ef fect is enhanced where ne w stars are being formed rather than in those zones where dust is destro yed by UV radiation. The chemistry of N 2 H + is also link ed to CO, and this molecule is particularly interesting, because it has been observ ed that its ab undance also tends to increase as the gas condenses to form stars, possibly in time scales of 10 5 yr This w ay the presence of both CS and N 2 H + molecules is a good indication of a possible ne w episode of formation. Ho we v er CO is a major destro yer of N 2 H + as the rst one e v aporates, and thus it will not be observ ed near re gions where ice is e v aporating from dust grains. Ho we v er if CO depletes, then the H +3 ab undance increases and N 2 H + is enhanced. This is kno wn to occur commonly in cold en v elopes. Could these observ ed v ariations in the location of the centers of emission a v ery rough astronomical clock? A v ery embedded cluster lik e cluster PL06 (Figure 5–6 ) sho ws a v ery clear coincidence between the peak of molecular gas emission and the location of the cluster itself. In the case of the partially emer ged cluster PL03, the gas emission contours sho w clearly the ef fect of ”hatching” in some of the most e v olv ed PMS stars. In

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121 the case of clusters PL04 and PL05, the ef fect reaches a possible e xtreme, as the emission of N 2 H + and C 18 O is limited only to those re gions in the cluster where remaining gas is dense enough to form ne w stars, and is practically undetectable in the re gions where cluster stars are already emer ging. Also we observ ed that a fraction of the stars in the clusters that locate outside the main gas emission re gions present e vidence of infrared e xcess, which suggests that the emer ging timescales of clusters could be similar or e v en shorter than the T T auri phase. 5.6 Summary and Discussion This chapter described our millimeter w a v elength observ ations of dense gas en v elopes associated with 8 embedded clusters in the Rosette Molecular Cloud. W e were able to construct high resolution inte grated emission maps for 6 molecular tracers to w ards the clusters. These maps sho w a clear relation between the locations of cluster stars and the remnant core gas. In most of the cases the clusters appear to be still deeply embedded in relati v ely compact gas clumps, although some of the clusters might be under going partial e v aporation as part of the early e v olution of the cluster W e estimated a series of basic physical properties for the gas en v elopes, and noticed that the most basic characteristics of the cluster en v elopes, namely size, mass and v elocity dispersion, are relati v ely uniform through the cloud. W e used sizes and v elocity dispersions to calculate virial masses, as well as total inte grated intensities to calculate column densities for all the tracers. T racer column densities were con v erted into H 2 column densities with the aid of A V estimates from the near -infrared colors of background stars in the observ ed areas. From the H 2 columns, we calculated then, the L TE gas mass totals in each cluster area. W e calculated ab undance ratios N ( X ) = N ( H 2 ) for each of the clumps in all of the tracers observ ed, and noticed that our v alues are reasonably similar those obtained in other studies that focus on indi vidual proto-stellar or starless cores. W e nd slight discrepancies in the a v erage ab undances with respect to the v alues referenced, b ut

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122 lar ge uncertainties in our v alues mak e it dif cult to point out for no w an y signicant dif ferences between the chemistry of cores associated with the formation of clusters, lik e in the Rosette, and those cores that present lo w or insignicant formation. F or all of the tracers the estimated virial to gas mass ratios suggest a lar ge percentage of gra vitationally unbound en v elopes, suggesting that the nuclei of the cloud cores where the clusters are located are in most of the cases out of equilibrium and possibly being held together only by an ambient pressure which we estimate to be of the order of 10 5 to 10 6 K cm 3 The virial mass M vir is quite sensiti v e to the line width of the clump spectra and the clump size. Because of our v ery localized clump area co v erage –which for some of our tracers might not be total–, the v alues could be systematically uncertain, b ut are in most cases lo wer limits. This w ay it might be possible that the virial to L TE mass ratios are indeed lar ge, and so RMC clumps could be already parting from virial equilibrium. As can be noticed from observ able v elocity gradients associated with the high density tracers (HCO + (1-0), CS(2-1) N 2 H + (1-0)), the RMC cluster forming clumps are still coherent in their most dense areas. On the other hand, lo w density gas, traced better by CO isotopes present small or ne gligible gradients, e vidence of more turb ulent edges. W e calculated star formation ef ciencies (SFE) for the indi vidual clusters by combining the estimates of total mass of gas and stars from our millimeter and near infrared w a v elength observ ations respecti v ely The ef ciencies of Rosette Clusters are belo w 10% for 7 of the clusters, while for the other tw o the ef cenc y is belo w 16%. This is in good agreement with clusters observ ed in other clouds. Our ef ciencies, ho we v er were calculated with only those cluster stars located ”inside” the projected area of the en v elope boundary In some of the clusters, a fraction of the stars might be already free of ambient gas b ut for those it is not possible an ymore to estimate the ef cenc y of formation.

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123 One standout characteristic of our observ ations is the e xistence of measurable of fsets between the peak of stellar density and the peaks of molecular emission for the dif ferent tracers. W e noticed that these of fsets were positi v ely correlated with the size of the cluster The lar gest clusters are also the most dense and the most populated, and also ha v e the lar gest formation ef cencies. W e noticed that the size of the clusters is not necessarily related to the size of the en v elopes, b ut the cluster to en v elope size ratios appear to be higher near the core of the cloud, where the most massi v e clumps are located. If those clusters at a more adv anced stage of emer gence are older than those which are still deeply embedded, then the relations we found are suggesti v e of a scenario in which the RMC started forming clusters at the cloud “Central Core”. This rst phase of formation w ould ha v e a v ery lar ge ef cenc y and resulted in a v ery rapid dispersion of the stars, with a timescale comparable to the T T auri phase of the cluster stars. According to Lada & Lada ( 2003 ), one characteristic of embedded cluster populations in GMCs is their v ery high inf ant mortality rate: most of embedded clusters will not survi v e to be bound systems. The main gra vitational binder of a young cluster is the gas itself, which pro vides the system with a deep potential well. If stars disrupt their molecular en v elopes rapidly the gas potential is destro yed before the system can form enough stars to a v oid f ast dynamical e v aporation. If we calculate the e xpected v elocity dispersions of the RMC embedded clusters as s cl = ( G [ M ? + M gas ] = R cl um p ) 0 : 5 and then the cluster crossing times t cr oss = 2 R cl um ps = s cl we can see that t cr oss v aries roughly from 0.9 to 1.2 Myr In order to remain bound, the time scales for the emer gence of the RMC clusters w ould need to be signicantly lar ger than these crossing times, which is unlik ely Considering the highly e xcited en vironment of the RMC, where the shock front from the neb ular ionization appears to promote rapid condensation of the gas near the center of the cloud, it is not dif cult to hypothesize rapid cluster formation follo wed by rapid gas e v aporation, especially in the presence of massi v e stars which could enhance disruption of gas en v elopes.

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124 It is v ery plausible that the clusters in the Rosette will become unbound stellar structures. As proposed by W ilking & Lada ( 1983 ), only clusters with ef ciencies closer or higher than 50% could survi v e as bound entities, which is v ery unlik ely to occur in the Rosette. These clusters will most probably increase the statistics of lar ge inf ant mortality rate in galactic stellar clusters. Finally we noticed that there were some observ able ef fects that relate the chemistry of the ambient gas and the formation of ne w stars in the cluster areas: in one cluster PL03, there is a visible separation between dif ferent tracer contour peaks, with CS(2-1) and N 2 H + (1-0) located near a group of highly reddened stars, and HCO + (1-0) apparently breaking into sub-clumps that follo w closely the morphology of near infrared bright neb ulosities – usually associated with reection by dusty material. Also, clusters PL04 and PL05 appear to be almost de v oid of gas, e xcept in localized areas where dense gas emission is signicant b ut rather weak if compared with highly embedded clusters lik e PL06 or PL07. Those areas are also populated with highly reddened objects. A tentati v e hypothesis is that once stellar winds remo v e material from the cluster centers, it condenses in other areas where ne w stars form. This ef fect has already been observ ed in other young clusters ( Langer et al. 1996 ; Whitmore 2003 ).

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CHAPTER 6 GLOB AL ASPECTS 6.1 A Near Infrar ed Extinction Map f or the RMC 6.1.1 Moti v ations Another direct application of our near infrared database is to map the distrib ution of dust through the cloud. In this chapter we tak e adv antage of the depth of our observ ations to construct an e xtinction map for the Rosette comple x from the near -infrared colors of the stars, using the techniques described by Lada et al. ( 1994 ) and Lombardi & Alv es ( 2001 ). This project has se v eral goals: a) to determine in detail the distrib ution of e xtinction for the cloud and within indi vidual cluster areas, b) to compare the characteristics of dark cores associated or not with star formation, c) to compare the properties of cores with those of the en v elopes studied in Chapter 5 and d) to determine an y possible relations between the properties of cores and their associated young clusters. The moti v ation behind these goals is that young clusters in the Rosette Comple x sho w di v erse stages of ”embeddeness”, as suggested by our observ ations of the local dense gas emission: we sa w that on the one hand, some clusters lik e PL01 or PL06, are clearly deeply embedded in thick en v elopes with a relati v ely uniform morphology On the other hand, clusters PL04 and PL05 seem to be at a later stage where a signicant fraction of the stars are already e xposed. In the e xtreme case, the clusters located in the Neb ula, particularly NGC 2244, are almost de v oid of parental dust and gas. Those clusters at partial or total e xposing stages in the neb ula and the cloud central core are, most probably slightly older than those which are deeply embedded, and thus a v ery interesting problem is to analyze the properties of e xtinction near these v arious cluster formation sites and obtain more information about the e v olution of young populations and their interaction with the local en vironment. 125

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126 Another independent moti v ation for this study is the lack of good quality e xtinction maps for the Rosette. The fe w that are a v ailable, are part of lar ge scale e xtinction studies, lik e those from Dobashi et al. ( 2005 ) or Cambr esy ( 1999 ). Ho we v er these maps are constructed from visual photometry counts (e.g. Digitized Sk y Surv e y), which are not capable of detecting background stars in re gions of high density in molecular clouds. V isual source counts are v ery limited in dynamic range, and are dominated by statistical uncertainties abo v e only 3-4 mag of visual e xtinction. Therefore optical counts cannot gi v e reliable estimates of the background density in re gions of acti v e star formation, where e xtinction can reach le v els of tens of visual magnitudes. Also, if the clouds are located at signicantly lar ge distances, lik e the Rosette, the counts of background stars are biased to w ards v ery bright stars, which also results in signicant uncertainties. Instead, Near -infrared w a v elengths (1.2-2.2 m) are 3 to 10 times less af fected by e xtinction ( Riek e & Lebofsk y 1985 ; Cohen et al. 1981 ) and thus ha v e dynamic ranges 5 to 10 times lar ger than visual w a v elengths ( Lada et al. 1994 ). Our surv e y although limited in co v erage, is sensiti v e enough to obtain reliable e xtinction estimates in re gions with e xtinction le v els of 30 visual magnitudes (e xtinction in the K band is approximately 10 times lo wer than in V ), accounting much better for background populations. These characteristics allo w for the construction of a much more detailed and reliable e xtinction map. 6.1.2 Dust Extinction fr om Near -Infrar ed Colors: NICE and NICER Originally Lada et al. ( 1994 ) presented the Near Infrared Color Excess technique (NICE), in which the e xtinction A V is estimated directly from H and K photometry This is done by calculating the dif ferences between colors ( H K ) of stars in a particular re gion with respect to the mean color of stars in a nearby control eld with minimum e xtinction, ( H K ) c This dif ference is kno wn as a color e xcess, E ( H K ) Using a normal reddening la w (e.g. Riek e & Lebofsk y ( 1985 )), the e xtinction to w ards a group of N embedded stars can be calculated simply as:

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127 < A V > = S R Ni = 1 [( H K ) i ( H K ) c ] N = S R < E ( H K ) >; (6.1) where the constant S R is equal to 15.87 for this particular reddening la w .. The more recently de v eloped Near Infrared Color Excess Re visited technique or NICER ( Lombardi & Alv es 2001 ), is an optimized e xtrapolation of the method just described. The main dif ference is that NICER mak es use of two independent colors, J H and H K and the normal reddening la w to estimate accurate v alues of the e xtinction A ( i ) V to w ards a group of N stars ( i = 1 ; :::; N ) The optimal e xtinction v alues are calculated by means of a tw o-dimensional maximum lik elihood scheme: by using the colors of stars in the control eld, one can estimate with a high precision the intrinsic scatter of the colors in a re gion with similar background eld density b ut v ery lo w e xtinction; then, by measuring the photometric scatter in the cloud eld, one has enough information to mak e a minimum v ariance estimate of the e xtinction. The e xtinction v alues can be calculated for indi vidual stars in the observ ed re gion and a v eraged within indi vidual pix els to create e xtinction maps with a certain resolution and signal to noise ratio. F or a group of N stars located inside one of these pix els, with center at position q = ( a ; d ) one can calculate the smoothed v alue of the e xtinction within that element as A v ( q ) = Ni = 1 W ( i ) A ( i ) V Ni = 1 W ( i ) ; (6.2) where W ( i ) is a weight v alue for the i th, star gi v en by W ( i ) = W ( q q ( i ) ) V ar ( A ( i ) V ) ; (6.3) where W ( q ) is a proper weight function. Using the method of Lombardi & Alv es we estimated the e xtinction v alues stars in the surv e y with J H and K bands a v ailable do wn to our completeness limit,

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128 Figure 6–1: NICER e xtinction map for the Rosette Comple x constructed with colors of stars from our surv e y catalog. The map w as constructed with a Nyquist sampling resolution of 27 arcsec, and smoothed with a Gaussian lter of FWHM=60 arcsec. The half tone le v els indicate steps of A V =1.0 mag, and the solid line contours indicate steps of A V =5.0 mag. K = 17 : 25 mag. W e discriminated stars with color errors lar ger than 0.1 mag (see chapter 4 ), and for this particular project we only used stars that f all inside the reddening band of the ZAMS, lea ving out infrared e xcess stars which ha v e an intrinsic red color due

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129 Figure 6–2: An inte grated emission 13 CO map of the Rosette Comple x by He yer et al. ( 2005 ) at the same scale as the e xtinction map of gure 6–1 The beam size of this map is 50 00 Half tone and solid contours indicate steps of 2 sigma abo v e the mean le v el. to circumstellar emission, as well as stars located to the left of the reddening band for which too blue H K colors are usually the result of high color dispersion and can be considered anomalous. W e measured e xtinction across the surv e y areas using a square resolution element of 54 00 which results in a nal Nyquist Sampled resolution of 27 00 or

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130 0 : 2 pc). In each box we estimated the v alues of A V for each star and stored the position of the star with the maximum e xtinction v alue; then we calculated the projected distance of e v ery other star inside the element to this e xtinction ”center”. The median v alue of these distances selected as the local angular size s W used in a Gaussian weight function of the form: W ( q q 0 ) = 1 2 prs 2 W e xp ( k q q 0 k 2 2 s 2 W ) ; (6.4) as suggested by Lombardi & Bertin ( 2001 ), and where r is the surf ace density of stars in the box ( n ( s t ar s ) = ar ea ( box ) ). F or those stars with no J band photometry a v ailable, we estimated A V ( i ) using the NICE method, described by equation 6.1. W e discarded those stars with a v ailable J and H b ut no K photometry because their color dispersion s J H is systematically lar ger than stars with good photometry in all bands. W e also discarded those stars with only one band photometry a v ailable, as for those it is not possible to obtain color information. The number of surv e y stars that were used to construct the e xtinction map is 78678, or approximately 54% of the total number of stars in our nal catalog. 6.1.3 The Extinction Map A smoothed map w as obtained by con v olving the resultant 27 00 resolution map with a second Gaussian lter of FWHM=90 00 In Figure 6–1 we present the e xtinction map, which is a good rendition of the distrib ution of dust e xtinction in the Rosette Comple x. White areas inside surv e y areas indicate zones with no photometric information a v ailable, which are a result of the zero point correction and discrimination applied to the data, as described in Chapter 3 W e nd an e xcellent agreement between this map and the distrib ution of 13 CO, mapped by He yer et al. ( 2005 ), which we sho w in Figure 6–2 for comparison. In gure 6–4 we compare the A V ( q ) with the logarithm of the 13 CO inte grated emission for each coincident line of sight in both maps. W ithin the range A V ( q ) = 0-10 mag, e xtinction

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131 Figure 6–3: This v ersion of the e xtinction map is equi v alent to the one of Figure 6–1 b ut for clarity halftone le v els ha v e been remo v ed, and contours are sho wn from A V =7.5 mag with steps of 0.25 mag. The starred symbols and labels indicate the locations of young clusters. is well correlated to the inte grated emission, with an inde x of 0.57. The least square t sho wn in the gure for A V =0 to 10 has a slope of 0.17. This correlation means that the re gions of high dust e xtinction are still well coincident with the re gions of highest density of molecular gas, just as in young molecular clouds, in spite of the onset of cluster

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132 Figure 6–4: Correlation between 13 CO inte grated intensity and A V at coincident positions within the surv e y areas. formation. The highest e xtinction re gions are in f act, still coincident with the zones of most acti v e star formation, which conrms the deep embedding and therefore the young age of the clusters. In gure 6–3 we present another v ersion of the e xtinction map that sho ws clearly the position of major clusters within the molecular cloud areas traced by dust e xtinction. W e were able to measure indi vidual e xtinction v alues do wn to a maximum of approximately A V = 58 mag for each star in the catalog, and to estimate mean e xtinction

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133 Figure 6–5: Distrib ution of indi vidual A V v alues between 0 and 30 mag. The thick solid line represents a linear t to v alues between 0 and 25 mag and is equi v alent to a po wer la w with e xponent -0.11. v alues in the range A V =0.0 to 32.0 mag per pix el across the surv e y areas. The e xtinction map has a pix el to pix el sensiti vity of 0.12 mag at a 1 sigma le v el. The map could be constructed with e v en a higher resolution, b ut the cost is the reduction of the signal to noise v alues and the creation of more “holes” in pix els with no photometry information. In gure 6–5 we sho w the distrib ution of indi vidual A V v alues in the range 0-30

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134 mag, which is monotonically decreasing and adjusts nicely to a linear t of the form l o g ( N ) = 0 : 11 A V + 4 : 48. It is f air to mention that the deniti v e correlation between e xtinction v alues obtained from near infrared maps and those inferred from the distrib ution of CO from radio observ ations, is still under study and a correct approach to the problem w ould require of careful statistical managing (see Lombardi ( 2005 )) which is be yond the basic scope of the thesis. Some current observ ational programs that aim to ha v e a complete understanding of the relation between dust and gas in star forming re gions in the conte xt of e xtinction ha v e already been successful in the study of some other star forming re gions ( Goodman 2004 ; Ridge et al. 2005 ) b ut a general scheme is still under de v elopment. It is interesting to notice ho w our map sho ws that e xtinction can be as high as A V = 5 8 mag outside the high density molecular gas re gions. The re gion of the Rosette Neb ula, around clusters NGC 2244 and NGC 2237 seem to be, as e xpected, mostly sweeped up of dust, although it is still permeated by a lo w uniform e xtinction layer This sho ws ho w the HI cloud still contains considerable amounts of dust, that ha v e not been destro yed yet by ultra violet radiation from the central OB association, as discussed in section 2.3.2 In general, a uniform 1-3 mag e xtinction seems to permeate the whole surv e y areas, with higher v alues delineating clearly the e xtension of the molecular cloud. 6.2 Indi vidual Extinction Cor es 6.2.1 Identication and Estimation of Pr operties W e isolated the areas corresponding to e xtinction cor es (local re gions of high e xtinction) associated with clusters, as well as those of “starless” cores, which we dene as non associated with cluster formation. W e dened a core as a re gion with at least one closed contour at a le v el equal or higher than A V = 10 mag. The local e xtinction maps for each one of these cores are sho wn in Figures 6–6 to 6–9 These maps ha v e the same resolution as the global map of gure 6–1 and the areas corresponding to the cores

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135 associated with clusters are equi v alent in size to the areas used for photometric analysis in Chapter 4 F or each one of the core areas we also constructed indi vidual e xtinction radial proles. These were done by measuring the mean v alues in concentric rings centered on the peak of e xtinction inside each indi vidual core area. These proles allo wed us to determine empirically rough plateau or “baseline” e xtinction le v el abo v e which the total mass of each core is calculated, from the proles we also dened core radii, which is dened as the point of intersection with these baselines. The mass of the cores were calculated by follo wing the reasoning of Dickman ( 1978a ) and T eix eira et al. ( 2005 ): Gi v en the v alue of e xtinction in each of the N indi vidual pix els of a map, the mass of molecular hydrogen associated with dust e xtinction can be estimated as M = ( a d ) 2 ( N H = A V ) N i A V ( i ) : (6.5) Here, a is the angular size of the resolution element in radians, d is the distance to the cloud in cm, N H = A V is the standard gas to e xtinction ratio of 1 : 9 10 21 cm 2 mag 1 ( Sa v age & Mathis 1979 ), where N H = N H I + 2 N H 2 and is the molecular weight corrected for Helium ab undance (which we assumed to be 0.28). Doing the corresponding con v ersions we obtain simply: M ( cor e ) = 2 : 59 N i ( A V ( i ) A V ( B )) M : (6.6) where A V ( B ) is the v alue of the baseline e xtinction considered for the particular core. In table 6–1 we sho w our estimations for core radii and mass in each one of the cluster forming and “starless” cores identied in the surv e y areas. W e include in the table the number of resolution elements in each map, the number of stars in the map that were “accepted” for the calculations of A V and the mean and maximum A V v alues per pix el.

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136 Figure 6–6: Indi vidual re gion A V maps for areas corresponding to clusters PL01 to PL06. Halftone le v els represent steps of 1.0 mag, and solid contours indicate steps of 5.0 mag. The plus symbols indicate the dened center of the corresponding A V core, and the cross symbol indicates the location of the cluster center

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137 Figure 6–7: Same as gure 6–6 b ut fort clusters PL07 to RLE10, NGC 2244, and NGC 2237.

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138 Figure 6–8: Same as gure 6–6 b ut for starless cores clusters SC1 to SC6.

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139 Figure 6–9: Same as gure 6–6 b ut for starless cores clusters SC7 to SC9. T able 6–1. Extinction Cores in the Rosette Comple x ID a RA b DEC b No. Stars c No. pix els d R cor e e M cor e f < A V ( q ) > s A V g max ( A V ( q )) h (J2000) [pc] [ M ] [mag] Cluster Host Cores PL01 97.9475 4.3375 750 81 0.9 423 8.02 2.95 16.73 PL02 98.3150 4.5850 839 100 0.9 519 5.51 2.99 15.84 PL03 98.3750 4.0300 1376 144 1.6 556 6.99 3.36 17.85 PL04 98.5425 4.3800 635 64 1.6 544 10.28 5.59 27.77 PL05 98.6225 4.2900 722 64 1.2 196 7.19 3.46 21.29 PL06 98.6525 4.2050 267 49 1.6 459 9.62 4.23 20.67 PL07 98.8700 3.9900 311 42 1.1 288 8.65 2.80 14.58 RLE08 98.5625 4.3300 482 64 1.4 850 12.13 5.80 28.04 RLE09 98.7700 3.7200 1059 144 1.9 736 8.98 5.40 32.00 RLE10 97.7575 5.2650 312 42 0.8 167 3.03 1.06 6.69 NGC2237 97.6250 4.9700 2076 210 1.9 670 3.14 1.83 14.04 NGC2244 98.0000 4.8600 2591 210 0.8 380 1.55 0.70 5.23 Starless Cores SC1 97.3975 4.9650 338 49 0.6 233 4.84 2.34 12.27 SC2 98.8500 4.0550 282 49 1.3 228 5.80 3.27 15.76 SC3 98.6100 4.0050 263 42 1.5 440 9.05 3.38 19.19 SC4 98.9075 3.7525 329 36 1.8 243 8.11 3.11 15.98 SC5 98.7600 3.8000 389 49 1.1 437 9.45 3.35 18.14 SC6 98.6700 4.6050 412 49 0.7 326 5.57 2.29 14.92 SC7 98.6725 4.1450 257 49 1.3 307 8.42 3.02 15.23 SC8 97.8375 4.3300 213 42 0.9 335 8.09 3.18 14.20 a Cluster host cores follo w nomenclature from Phelps and Lada (1997) and Rom an-Z u niga et al. (2006a). Starless cores are indicated with nomenclature ”SC”. b Core center dened as maximum A V ( q ) location. c Number of stars in selected core area d Number of pix el elements in selected core area e Core radius. f Core mass, from equation 6.5 g Mean and standard de viation of e xtinction per pix el in selected core area h Maximum e xtinction per pix el in selected core area

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140 Figure 6–10: Extinction proles for areas corresponding to clusters PL01 to PL06. The distance is measured from the center of cores, which are indicated in gures 6–6 to 6–9 Each radial step is equi v alent to 1/20th of the size of the corresponding core map.

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141 Figure 6–11: Same as gure 6–10 b ut fort clusters PL07 to RLE10, NGC 2244, and NGC 2237.

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142 Figure 6–12: Same as gure 6–10 b ut for starless cores clusters SC1 to SC6.

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143 Figure 6–13: Same as gure 6–10 b ut for starless cores clusters SC7 and SC8. 6.2.2 Cor e Sizes The distrib ution of core sizes obtained from the radial proles is sho wn in Figure 6–14 W e plotted separately the distrib ution for cores associated with clusters and starless cores. Cluster forming cores are lar ger in the a v erage, with a median of 1.54 pc, while the starless cores ha v e a median of only 0.98 pc. F or comparison, we included in the plot the equi v alent radii of clusters obtained from the analysis of Chapter 4 The cluster equi v alent radii R eq are distrib uted around 1.0–1.25 pc, which suggests that in most cases the clusters should be enclosed in their original cores. In Figure 6–15 we compare the equi v alent radii of clusters R eq with the size of their corresponding cores R cc W e see ho w 7 out of 12 clusters are either more e xtended or ha v e comparable sizes to those of their parental cores. This group of clusters include PL01, PL02, PL03, PL04 and the three clusters in the neb ula area, RLE10, NGC 2244 and NGC 2237. The remaining clusters, PL06, PL07, RLE08 and RLE09 are indeed the most hea vily reddened of the cloud, and are the only ones with a cluster to core ratio of less than 1.0. One particularly interesting case is PL05, for which the core is v ery localized b ut lar ge enough to still occup y a signicant fraction of the cluster area. Ho we v er part of this cluster is already emer ged, and the emission of dense gas, as we sa w in the last chapter is conned to the areas of hea viest oscuration, where v ery red

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144 Figure 6–14: Distrib ution of radii for starless cores, cluster host cores, and embedded clusters in top, mid and bottom panels respecti v ely stars suggest that at least the core of this cluster is stil under going formation (see also appendix C .) W e can e xplore this idea a little bit further by comparing the cluster equi v alent radii vs. the mean e xtinction v alues of their parental cores, as we do in the top panel of Figure 6–16 W e see ho w the most e xtended clusters are also the ones that present less obscuration. All of the cores associated with the main molecular cloud areas ha v e mean e xtinction v alues abo v e the 4.0 mag le v el, and their sizes are ne gati v ely correlated to

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145 Figure 6–15: One to one cluster vs. parental e xtinction core sizes are indicated by the points. The dashed line indicates a ratio equal to unity e xtinction (the correlation coef cient is h = 0 : 55). In contrast, the positi v e correlation ( h = 0 : 57) between e xtinction and IRX fraction sho wn in the bottom panel of the gure, suggests that those clusters emer ging from their parental cores are precisely the more e v olv ed. Because the IRXF are signicant, the parental core dispersion must occur quite rapidly possibly before the end of the circumstellar emission phase. Then we compared the mean e xtinction in clusters v ersus the sizes of the corresponding cores, as sho wn in Figure 6–17 There is also an observ able correlation between mean e xtinction and the size of cores, which indicates that those cores with lar ger column densities are also the ones with lar ger v olumes. Clearly the e xposed clusters in the neb ula area ( A V < 4 : 0mag) for which the cores are small and brok en, stand out from this correlation. 6.2.3 Cor e Masses In general, e xtinction cores associated with cluster formation are not only lar ger b ut also might be more massi v e, as sho wn in Figure 6–18 The mean v alue for masses of star forming cores M cc is 438 M v ersus only 319 M for the starless cores masses, M sc The distrib ution of M sc is v ery narro w and apparently v ery dif ferent from the M cc distrib ution, which is almost normally distrib uted. The number of cores in each group

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146 Figure 6–16: T op: cluster equi v alent radii vs. median e xtinction. The solid line and dotted lines represent a least squares linear tting 1 rms de viation. Bottom: IRX fraction ( K < 15 : 75 mag) vs. median e xtinction is v ery small, and thus it is not possible to determine if the distrib utions are complete at both ends. Ho we v er when summed together the distrib ution of core masses appears to ha v e a monotonic increase from the lo west mass end, and then an abrupt decrease in the lar gest mass bin that might suggest that there are only a limited number of v ery massi v e star forming cores in an acti v e cloud lik e the Rosette, which is consistent with the observ ations of W illiams et al. ( 1995 ).

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147 Figure 6–17: Mean e xtinction v ersus cluster forming core size. The v alues are correlated with an inde x of 0.53, and a least square t to the data is sho wn by the solid line. W e also w ould lik e to compare the mass of the cores with the mass of the cluster products themselv es. Cluster masses M cl us were determined with aid from the embedded clusters IMF models of Muench et al. ( 2000 ), follo wing a similar method to Lada & Lada ( 2003 ) (see section 6.2.4 ). Using the IMF models we were able to estimate a luminosity to mass con v ersion f actor for the total number of stars located within the cluster boundaries dened in Chapter 4 The counts included stars do wn to the completeness limit K < 17 : 25 mag., and were corrected by the contrib ution of the uniform background eld. The con v ersion f actor w as calculated for a cluster with an age of 1 Myr follo wing the PMS e v olution models of D'Antona & Mazzitelli ( 1997 ), and assuming an IMF similar to the T rapezium cluster at the distance of the Rosette and a distrib ution of e xtinction similar to the one observ ed in each of the indi vidual cluster areas. In the left panel of Figure 6–19 we compare the mass distrib ution for M cc and embedded clusters. The distrib ution for embedded clusters peak at a much lo wer v alue, which results in moderate star formation ef cencies. In the right panel we separated those clusters for which we ha v e IRAM 13 CO(2-1) data (PL01 to PL07 and RLE08) and

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148 Figure 6–18: Distrib ution of e xtinction core masses. The top tw o panels sho w separate distrib utions for cores associated with clusters and starless cores. compared their corresponding mass distrib ution with those from e xtinction cores and 13 CO(2-1) en v elopes in both virial and L TE equilibrium. As we can see, the distrib ution calculated for cluster local en v elopes is similar to the one from e xtinction cores. This sho ws a good consistenc y among our independent radio and near infrared observ ations. 6.2.4 The Embedded Cluster Mass Function In their recent re vie w Lada & Lada ( 2003 ) constructed a v ery complete catalog of nearby ( d < 2 kpc) embedded clusters from the a v ailable literature. F or each cluster

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149 Figure 6–19: Left: The shaded histogram is the mass distrib ution of embedded clusters, M cl us and the thick solid line histogram is the mass distrib ution of e xtinction cores. Right: Same as in the the left panel, b ut only for those clusters which ha v e IRAM counterpart observ ations. The lled circle symbols joint by a dotted line represent the virial mass distrib ution for cluster en v elopes, and the cross symbols connected by a dashed line represent the L TE mass distrib ution for cluster en v elopes, as discussed in chapter 5 in the catalog the y included a total number of stars detected, and from those, the y used the models of Muench et al. ( 2000 ) to estimate indi vidual cluster masses. Lada & Lada assumed that the IMF for all of the clusters in their catalog w as similar to the one for the T rapezium cluster Ho we v er the y used tw o dif ferent PMS e v olution models for the calculations of the luminosity to mass con v ersion f actor depending on the reported ages of the clusters in the list: one model is for a cluster of 0.8 Myr of age, similar to the T rapezium, and one of 2 Myr similar to the cluster IC 348 in the Perseus Molecular Cloud. The y also corrected the counts according to the distance to the corresponding molecular clouds, and added an a v erage e xtinction v alue of A V =5.0 mag. Using these estimations of the masses of embedded clusters in nearby clouds, the y constructed the Embedded Cluster Mass Distrib ution Function (ECMDF), by adding indi vidual cluster masses in logarithmic mass bins of 0.5 de x in width. The ECMDF constructed this w ay is the distrib ution of total cluster mass, N M cl us as a function of cluster mass in logarithmic units (M cl us ). The minimum M cl us in their function

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150 w as 35 M corresponding to the cluster MWC 137, and the maximum one, 1100 M corresponded to the ONC/T rapezium cluster In Figure 6–20 we present a v ersion of the same gure; just as sho wn in Lada & Lada ( 2003 ) we also included in the plot the mass distrib ution functions of clusters from the catalog of Hodapp ( 1994 ) and the one constructed with those clusters located within 500 pc of the Sun, which are the best studied. Figure 6–20: The Embedded Cluster Mass distrib ution function, for the catalog of Lada & Lada ( 2003 ). The solid histogram represents the mass distrib ution of all clusters in their catalog, while the dotted line function includes only those from the outo w sample of Hodapp ( 1994 ). The dashed line function represents only those clusters within 500 pc of the Sun (distance limited (D.C.) sample). Reproduced with permission from the authors.

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151 Using the same method, we constructed then a ECMDF for the clusters in the Rosette and also one which separates the contrib ution for the clusters of the Orion Molecular Cloud in the database of Lada & Lada The resultant distrib utions, compared to the ECMDF are sho wn in Figure 6–21 Figure 6–21: Comparison of the ECMDF of Lada & Lada ( 2003 ) (dashed line) with equi v alent distrib utions constructed with clusters from the Rosette (solid thick line) and Orion (dotted line). The cross symbols sho w the total ECMDF after adding the Rosette clusters. T w o important aspects can be highlighted from this gure: The rst one is that our deri v ed ECMDF for the Rosette is consistent with the one for the Orion Comple x

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152 within the range (1 : 25 < log M cl us < 2 : 25 M ), which contains the lar gest number of clusters. This can be interpreted as a similar mass spectral inde x, d N = d M for clusters masses of 20 to 180 M where the lar gest number of clusters is contained (5 Orion clusters and 7 Rosette clusters). In the third bin (2 : 25 < log M cl us < 2 : 75 M ), the Rosette contrib utes three clusters and Orion with clusters, while the last bin (2 : 75 < log M cl us < 3 : 25) is populated by only one Orion cluster: T rapezium ONC. This means that the contrib ution to the mass distrib ution function in the upper end of the spectrum comes basically from a v ery fe w lar ge clusters. When the contrib ution of the RMC clusters is added to the ECMDF (cross symbols), the function stays the same within the errors, which conrms one important nd in the study of Lada & Lada : the ECMDF is roughly at between 50 to 1000 M A at mass distrib ution corresponds to a cluster mass spectrum with an spectral inde x of -2, which is similar to the spectral inde x -1.7 of dense molecular cloud cores ( Lada et al. ( 1991a )), and could suggest a uniform star formation ef cenc y for most cluster forming cores. The second aspect is that the drop of f in the contrib ution of the rst bin (1 : 25 < log M cl us < 1 : 75 M ) is signicant in the Rosette and Orion as it is for the whole database, which conrms that lo w mass clusters are rather scarce entities, and that in most of the cases, stars are formed in rich clusters, with masses of at least 50 M with no more than 10% of the stars in the galaxy being formed in small groups (see also Porras et al. ( 2003 )). In our surv e y almost 86% of the total embedded stellar mass is contained in the lar gest 9 clusters, which all ha v e masses lar ger than 50 M b ut we also found three clusters with masses between 20 and 50 M as well as a small contrib ution from a lo wer density population, which, if it w as not for the high contrib ution of the eld, maybe could be separated into small groups that w ould sum up to a contrib ution comparable to the lo w mass bin counts of Hodapp ( 1994 ) (see Figure 6–20 ). In f act, the non-detection of small groups comes from our denition of cluster in which we require a

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153 minimum number ( j = 10) of IRX neighbors; an y group with less members will simply not be detected and thus not be considered here. 6.2.5 Star F ormation Efcencies Star F ormation Ef ciencies SFE, were calculated with respect to the star forming cores as M cl us = ( M cl us + M cc ) and are listed in table 6–2 These total SFEs are dif ferent from the local SFEs calculated with respect to the molecular gas en v elope masses, as those were estimated with the embedded cluster masses, M emb only and ga v e an estimate of the production of stars in the specic re gions of high density gas tracer emission. When we use M cl us and M cc there are some cases for which M cl us > M cc which underestimates the original core mass that went into stellar production. This results in SFEs abo v e 50%, which can only be considered as upper quotas. In Figure 6–22 we sho w an interesting correlation between the total SFE from table 6–2 and the equi v alent radii of the corresponding cluster which sho ws that those clusters with higher ef ciencies are also the most e xtended. W e discarded for this plot the embedded cluster PL05 and the clusters in the neb ula area, for which the cores are almost completely dispersed and the SFE v alues are uncertain. The correlation inde x for the 8 pairs used in the gure is 0.72. This correlation might suggest that the higher ef ciencies present in the lar ger more dispersed clusters are e vidence of this clusters reaching the latest stages of their embedded phases, and this is done after reaching an ef cienc y limit. Ho we v er as discussed in chapter 5 the RMC is a cloud with a lot of stimulation due to the presence of the neb ular shock front, and this stimulation might promote a rapid condensation of the molecular cores and a v ery ef cient production of clusters, which causes the rapid dispersion of the core remnants. The prime consequence of this ef fect is a v ery lo w probability for the clusters to remain gra vitationally bound, because the y are remo ving their gas v ery f ast. In general, total SFE v alues are lar ger than the local SFE v alues calculated in chapter 5 because in that case we were only considering those stars that were located inside

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154 Figure 6–22: SFE calculated for the e xtinction cores vs. equi v alent radii of embedded clusters. Cores with ef ciencies abo v e 50% were discarded (see te xt). The solid line is a least square t to the data. The pairs are correlated with an inde x of 0.72. the observ ed en v elopes. The main result from that chapter is that locally the ef cienc y of formation is lo w No w that we consider the e xtinction cores in the conte xt of complete clusters, we see that in tw o cases, PL05 and NGC 2244, clusters are signicantly more massi v e than the cores, resulting in ef ciencies close to 50%. This, ho we v er could suggest that there is a maximum ef cenc y of formation for embedded clusters before the y start to get rid of a signicant amount of gas. Other clusters lik e PL03 and PL04, which might be at slightly earlier stages of dispersion, with ef ciencies between 20-30%, and if we consider them clusters that are close to a point of “no return” where catastrophic dispersion of gas is ine vitable, then 20-30% could in principle be a good estimate of a maximum possible ef cienc y prior to gas dispersion in embedded clusters. Numerical studies suggest similar numbers: for e xample, Lada et al. ( 1984 ) determined that bound clusters could form with ef ciencies as lo w as 30% if gas disperses slo wly o v er time scales of about 3.0 Myr b ut in principle ef ciencies of 50% or more are required if gas e xpulsion is quick. Under this picture, if the spectroscopically determined age 1.9 Myr NGC 2244 is correct and the clusters in the molecular cloud are

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155 Figure 6–23: V ariation of SFE as function of distance to NGC 2244. The error bars for clusters with ef ciencies belo w 35% are calculated as the standard de viation of v alues. Higher ef ciencies are considered as upper quotas because the gas mass might be underestimated. The dashed v ertical lines at 10, 20 and 30 pc indicate roughly the edges of the neb ula central ca vity the Monoceros Ridge and the cloud Central Core, respecti v ely younger then it is v ery unlik ely that these clusters will survi v e as bound entities, because we already see signicant gas dispersion. Ho we v er more recent numerical studies with a more precise treatment of stellar encounters, lik e Kroupa et al. ( 2001 ) suggest that ef ciencies of 30% could be enough to assure gra vitational survi v al for clusters e v en in the presence of rapid gas e xpulsion. Also, Boily & Kroupa ( 2003 ) found that if the stellar v elocity distrib ution in a cluster f a v ors stars with lo w v elocities, then a system in virial equilibrium could survi v e disruption e v en with an ef cienc y of less than 40%. Finally in Figure 6–23 we present the v ariation of SFE as a function of the projected distance to the center of NGC 2244 which coincides with the Rosette Neb ula. In the gure we also indicate marks at 10, 20 and 30 pc, which roughly coincide with the projected edges of photoionization (central ca vity), the Monoceros Ridge and the projected edge of the Neb ula as indicated by Celnik ( 1986 ). F or clarity we represent

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156 Figure 6–24: This map of the Rosette Comple x sho ws the conte xt of the interaction between the Rosette Neb ula –sho wn by halftone le v els of 25 m emission– and the Rosette Molecular Cloud –sho wn in contours of 13 CO emission–. The dashed circles at 10, 20 and 30 pc indicate roughly the edges of the neb ula central ca vity the Monoceros Ridge and the cloud Central Core, respecti v ely The locations of the clusters are indicated by lled circle symbols. these edges schematically in the map of Figure 6–24 The formation ef ciencies peak at the cloud Central Core, located between 20 and 30 pc from NGC 2244, inside the neb ular edge. The clusters PL07 and RLE09 located outside of ionization edge, as well as the clusters PL01 and PL02, RLE10 and NGC 2237, located at the edge of the central ca vity

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157 T able 6–2. Star F ormation Ef ciencies ( A V cores) Cluster SFE [%] PL01 12.6 PL02 10.0 PL03 21.8 PL04 30.8 PL05 < 50.9 a PL06 9.0 PL07 6.9 RLE08 9.8 RLE09 6.8 RLE10 18.6 NGC 2237 20.4 NGC 2244 < 51.8 a a F or these clusters M cl us > M cc and the gas mass is underestimated. sho w signicantly smaller ef ciencies. These dif ferences, as we discuss belo w are the result of dif ferent le v els of cluster formation stimulation across the cloud, directly related to the cloud-neb ula interaction. 6.3 Summary and Discussion In this chapter we discussed some global aspects of the relation between the Rosette Comple x clusters and their formation en vironments. The main tool for this brief analysis w as an e xtinction map constructed with the near -infrared colors of our FLAMINGOS database. Our map allo ws to trace re gions with e xtinction le v els of up to 30 mag., almost 10 times deeper than studies based on visual photometry counts. The distrib ution of e xtinction across the Rosette comple x observ ed in this map is well correlated with CO emission traced with sub-millimeter observ ations. W e were able to detect in our map signicant e xtinction cores associated with each one of the 9 clusters identied in Chapter 4 as well as 8 other cores not associated with cluster formation, or “starless”. W e calculated and analyzed the distrib ution of sizes and masses for these cores, and compared them to the sizes and masses of the clusters in the comple x. W e found that those cores associated with cluster formation are lar ger and more massi v e in the a v erage, which is consistent with the results obtained from analysis of the molecular clumps ( W illiams et al. 1995 ).

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158 The general result obtained from this analysis is that the amount of e xtinction present to w ards a cluster is correlated ne gati v ely with respect to its size ( R eq ), b ut correlated positi v ely to its IRX fraction. This suggests that emer gence of embedded clusters from their parental cores might be follo wed by a relati v ely rapid dispersion of the stars: W e found that 7 out of 12 clusters in the Rosette are lar ger or close in e xtension to their associated e xtinction cores, which suggests that most of the clusters are old enough to suf fer detectable parental material remo v al b ut, because the clusters ha v e signicant IRX fractions, then the time scales for gas dispersion must be shorter or at least comparable to the T T auri phase, which is also consistent with the observ ed morphology of the dense gas emission discussed in chapter 5 W e also included in this chapter a rst approximation to the calculation of total stellar mass for the Rosette clusters, using a constant luminosity to mass con v ersion f actor calculated with the models of Muench et al. ( 2000 ). The distrib ution of cluster masses obtained this w ay is consistent with the one obtained for clusters compiled in the catalog of Porras et al. ( 2003 ), and conrms that the embedded cluster mass distrib ution is approximately at between 50 and 1000 M which results in a similar mass spectral inde x to that one of molecular gas clumps, suggesti v e of a uniform star formation ef cienc y in GMCs. The cluster mass distrib ution is also similar to the one of the Orion Comple x, and we found that for both clouds, the drop in the number of clusters with observ ed masses of less than 50 M is signicant. This suggests the e xistence of a minimum cluster mass that is a result of the minimum number of stars that can be considered to form a cluster: according to dynamical estimates by Adams & Myers ( 2001 ), this number should be around 35 stars, which is in f act close to the number of members we observ e in clusters PL06, PL07 and RLE10, the smallest we identied in the comple x. Belo w this minimum number and mass, dynamical survi v al of clusters is not assured and thus, clusters smaller than this limit may not e xist, b ut here we need to tak e into account that our detections are also articially constraint: F or e xample, the nearest

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159 neighbor technique we discussed in chapter 4 w as not parameterized for the detection of groups with 10 IRX members or less. Ho we v er such small stellar groups w ould need to be considered as products of a dif ferent formation process. Our Rosette cluster mass estimates allo wed us to calculate star formation ef ciencies (SFE), which we found to be positi v ely correlated to the equi v alent sizes of the clusters, which may suggest that lar ger ef ciencies resulted in rapidly e xpanding clusters. The SFE, ho we v er can only be correctly estimated for clusters with a total stellar mass smaller than their core mass, and for those clusters with positi v e M cl us = M cc ratios, the resultant SFE can only be considered as an upper limit. The SFE appears to be signicantly lar ger for clusters located between 20 and 30 pc from NGC 2244, which corresponds to the cloud central core. He yer et al. ( 2005 ) sho wed that those re gions of the cloud located inside the 30 pc edge, namely the cloud Central Core and the Monoceros Ridge, ha v e signicantly lar ger turb ulence scale e xponents, as well as a more coherent v elocity structure, which are suggesti v e of a direct inuence from the e xpanding neb ula in the cloud dynamics behind the neb ular e xpansion front. This interaction, ho we v er may not be a unique trigger for cluster formation in the RMC, because there we observ e tw o clusters, PL07 and RLE09, in the cloud Back Core, located be yond the ionization front. These tw o clusters ha v e smaller ef ciencies, comparable to those of Clusters PL01, PL02 located between 10 and 20 pc from NGC 2244. This is interesting, because as we discussed in chapter 5 the dense gas emission contours of cluster PL02 sho w strong gas compression in a direction perpendicular to the HII sphere e xpansion, which may suggest that the Ridge clusters may suf fer from strong clump photoe v aporation by the winds of NGC 2244. In other w ords, for these clusters the interaction with the ionization front may actually be a ne gati v e inuence in the formation ef cienc y What this suggests is that the inuence of the OB association is not uniform across the cloud, and stimulated formation can be in f act a relati v e ef fect. A correct interpretation of the ef fects of the interaction between the tw o parts of the comple x

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160 (neb ula and cloud) will require careful modeling of the gas dynamics and an accurate estimation of the time scales for cluster formation and local en vironment dispersion.

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CHAPTER 7 CONCLUSIONS AND FUTURE W ORK 7.1 Distrib ution of Y oung Stellar P opulations in the Rosette Complex As part of a lar ge research program dedicated to the in v estigation of the global aspects of star formation in GMCs we selected the Rosette Comple x, a prominent astrophysical laboratory located at 1.6 Kpc from the Sun in the constellation of Monoceros. Using FLAMINGOS, a po werful near -infrared imager we were able to surv e y the areas of the Rosette Comple x in J H and K do wn to a completeness limit of K = 17 : 25 mag, almost 3 magnitudes deeper than pre vious studies and enough to account for stellar masses do wn to the HBL. The quality of our data w as enough to obtain reliable photometric colors do wn to a limit of K = 15 : 75 mag, which we used to detect stars with signicant near -infrared e xcess, which trace the youngest stellar populations. The tracing of re gions of high surf ace density of IRX sources sho wn unequi v ocally the locations of the most recent episodes of formation, and allo wed us to analyze the distrib ution of stars across the comple x. The follo wing are the highlights of this analysis: Our observ ations conrmed the e xistence of se v en pre viously detected young clusters and re v ealed tw o more, which are e v en more deeply embedded. W e also were able to detect tw o ne w clusters in the area of the Rosette Neb ula, and to conrm a small b ut signicant number of young sources in the NGC 2244 association. The 9 embedded clusters in the Rosette Molecular cloud account for approximately 87% percent of the total young population, which is similar to what has been observ ed in other clouds. When we add to the counts the clusters located in the neb ula areas, the fraction of stars in clusters is reduced to 60% which could be the result of a lar ger contamination by eld IRX stars outside of the molecular cloud re gions and a signicantly smaller fraction of IRX sources in the neb ula clusters. 161

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162 W e found that up to 50% of the total number of sources in embedded clusters is contained in 4 clusters, which locate at the most dense re gion of the molecular cloud, kno wn as the Central Core. Three of the clusters in this re gion seem to be part of a v ery lar ge structure of cluster formation, and are signicantly e xtended and populous. The fourth cluster in the group is not as numerous b ut contains the most massi v e star detected in the cloud, a B binary kno wn as AFGL-961. 7.2 The Local En vir onments of Y oung Clusters W e also studied the interaction between the Rosette clusters and their local formation en vironments, with a complementary surv e y in millimeter w a v elengths for eight of the embedded clusters. The main results of this complementary project are: All of the clusters areas present signicant dense gas emission. The morphology of this emission suggests a tight correlation between the distrib ution of stars and gas, b ut also dif ferent stages of embeddeness across the cloud: on one hand, some clusters are associated with well dened, compact en v elopes that surround the entire cluster areas. On the other hand, some clusters sho w a more loose emission, v ery weak in what seem to be the e xposed parts of the clusters, and v ery intense at the areas of more recent formation, which are e videnced by highly reddened sources. In the latter case, clusters might be sho wing e vidence of partial emer gence from their parental clouds, and could be slightly older than those with well dened en v elopes. The peaks of stellar density and tracer emission presented in most cases signicant projected of fsets which are positi v ely correlated with the sizes of the clusters, conrming that the most e xposed clusters are also the most e xtended. Ho we v er in the a v erage, the basic properties of the dense gas clumps (size, mass and line width) are relati v ely similar among the sample, which suggests that locally the en vironment of clusters are similar to each other Approximately 50 to 75 percent of the cluster en v elopes sho w bound virial to L TE gas mass ratios when measured with CS(2-1) and HCO + (1-0), b ut this fraction is smaller

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163 when measured with N 2 H + (1-0) 13 CO(2-1) and C 18 O(2-1)(14, 33 and 38%) respecti v ely This appears to suggests that the hypothesis of local equilibrium might not hold at the re gions of the RMC that coincide with the clusters, and a possible e xplanation is the interaction of the dense gas with the onset winds of the young stars. W e found that if the en v elopes are not gra vitationally bound, then the y must be constrained by e xternal pressure from the medium, and this pressure w ould be of the order of k 10 5 km cm 2 Also, the v elocity gradients in the gas clumps are lar ger for high critical dense tracers lik e CS(2-1) and HCO + (1-0), than those measured with 13 CO(2-1) which is suggesti v e of more coherent gas motions at the most dense parts of the en v elopes, which w ould be surrounded by more turb ulent edges. As part of our analysis, we were able to measure ab undance ratios for the dif ferent tracers in each cluster area. The ratios we obtain are in v ery reasonable agreement with other studies a v ailable in literature, which ho we v er are mostly dedicated to starless molecular clumps. This study is among the fe w that report ab undances for cluster forming clumps, which is a v aluable addition to these in v estigations. 7.3 Extinction in the Rosette Complex and Global Results Using the near -infrared colors of sources in our surv e y we constructed an e xtinction map. This map has tw o main adv antages o v er other a v ailable maps for the Rosette re gion made from visual observ ations: a better resolution (90 00 after smoothing) and a much lar ger dynamic range (0 to 30 mag). This map allo wed us to determine the location and e xtensions of molecular cores associated with dust emission across the comple x. W e found signicant cores associated with the 9 embedded clusters, plus 8 more that are not associated with star formation. W e calculated sizes and masses of these cores which allo wed us to gi v e more insight to our analysis: The amount of e xtinction to w ards a cluster decreases in proportion to the e xpansion of the cluster as it emer ges from its parental core. Se v en out of 12 clusters in the Rosette ha v e sizes comparable or lar ger than their e xtinction cores, which suggests that the

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164 embedded clusters are old enough to under go signicant parental material remo v al. Ho we v er because the clusters still present a signicant fraction of members with infrared e xcess, the time scales for parental gas remo v al must be shorter than the T T auri phase. In other w ords, remo v al is relati v ely rapid after cluster formation. W e calculated masses for the embedded clusters themselv es, and observ ed the distrib ution of these masses with respect to the one observ ed for other clusters in the galaxy W e found that the mass distrib ution of clusters in the Rosette is similar to Orion, and that its contrib ution to the global cluster mass function of the galaxy does not alter the result of Lada & Lada ( 2003 ): the mass distrib ution for young clusters is at in the range of 50 to 1000 M which is suggesti v e of a relati v ely uniform star formation ef cenc y for dif ferent clouds in the galaxy By combining our cluster and core mass estimates we were able to calculate accurate star formation ef ciencies. Although for the tw o more e xposed clusters SFE v alues can be considered only as upper quotas, we found that a maximum ef cienc y of 20 to 30% could be achie v ed in a cloud lik e the Rosette. The ef cienc y of formation is clearly higher at the Central Core of the RMC, with v alues belo w the threshold at the Monoceros Ridge and the Back Core. In the rst case, the ef cienc y sees to be enhanced by the interaction of the e xpansion front of the neb ula and the molecular cloud; in the second case the ef cienc y could be ne gati v ely inuenced by photoionization from the winds of the massi v e stars in NGC 2444; in the third case, the inuence of the neb ula e xpansion is minimal, which results also in lo wer ef ciencies b ut at the same time discar ds sequential star formation as the unique trigger for cluster formation in the comple x. 7.4 Futur e W ork This w ork is the result of a lar ge process of learning, which started with the testing of the instrumentation we used, continued with a long term assessment of the data quality and the de v elopment of techniques for their analysis, and nished with an in v estigation of some of the scientic aspects that were rele v ant for our initial goals. Ho we v er our

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165 data still ha v e information to be e xtracted and other analysis needs to be done for a more complete understanding of the problem we originally laid out. I w ant to think of this as a solid base for a long term project, which could be b uilt with more observ ations, contrib utions from numerical models and nally a link to theoretical studies. F or e xample, we ha v e already observ ed 23 multi-object spectroscopic plates with FLAMINGOS, which contain information for a fe w hundred stars across the Rosette comple x. These data are still under analysis, as we ha v e to solv e some preliminary problems, lik e the construction of an accurate near infrared spectral classication sequence for stars of early types (O, B, A, F and G) –which are the ones we were able to observ e in the Rosette gi v en its distance. The FLAMINGOS team at the Uni v ersity of Florida already constructed successful sequences for late type stars, which can be classied independently of the resolution of spectral line equi v alent widths, b ut that is not the case for early types. Once this problem is solv ed, we should be able to classify bright members of the young population of the Rosette with good accurac y and this will allo w us to estimate mean ages for those embedded clusters with 20-30 observ ed members 1 The age estimation, ho we v er is e xpected to be some what challenging, as we w ould e xpect v ariations among the relati v e ages of the clusters of at most a fe w hundred thousand years. The main sources of error in the determination of ages of young stars w ould come from the quality of the data, b ut also there are intrinsic sources of error in the stars, lik e photometric v ariability and spectroscopic v eiling ( Hartmann 2001 ). Most probably impro v ed methods of stellar age-dating will require e xtensi v e Monte Carlo 1 Most probably clusters lik e PL01, PL02 or PL06 will ha v e less than this number of members observ ed, because their projected sizes are small compared to the co v erage of the MOS plate (10 arcmin length at the KPNO 4m telescope). Those clusters might need to be re-observ ed with a lar ger telescope lik e Gemini. The second generation instrument FLAMINGOS 2 will be required for those observ ations

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166 simulations of each one of the problematic parameters and repetition of some of the w orst quality observ ations, which still ha v e to be assessed. After the ages of the clusters are determined from the spectroscopic observ ations, we should be able to ha v e a v ery good estimate of the IMF of the clusters do wn to 1 M approximately This, along with carefully constructed luminosity functions for each cluster should be enough to parameterize the models of Muench et al. ( 2000 ) and determine accurate IMF for each cluster T w o v ery interesting problems to attack w ould be a) the uni v ersality of the IMF across the cloud b) the global ef cienc y of star formation for high vs. lo w mass stars. Additionally we should be able to calculate star formation rates and to determine ho w the y v aried across the cloud. This information w ould allo w us to reconstruct the process of formation in the Rosette Comple x and to compare to numerical models. The data from the Spitzer infrared space telescope seems to be the most reliable source of counterparts of our young stellar source candidates in the mid infrared re gimes which w ould conrm a) the circumstellar disk fractions of the Rosette clusters, b) the e xistence of circumstellar disk candidates not associated with clusters (distrib uted population). It has been sho wn that the 3.6-8 micron images obtained with the instruments IRA C and MIPS impro v e signicantly the accurac y of the determination of circumstellar disk candidates because of the possibility of obtaining direct measurements of the total disks emission from the spectral ener gy distrib utions of indi vidual stars (see e.g. Muench et al. 2005 ). A v ery signicant fraction of the area of the Rosette Molecular Cloud is already being imaged with the Spitzer telescope (Ian Bonnel et al.), and we are ne gotiating a collaboration for the joint analysis of our databases. Another database that w ould also gi v e reliable IRX counterparts for our near -infrared observ ations is the IPHAS surv e y( Dre w et al. 2005 ), which will obtain accurate imaging of all star formation re gions in the northern hemisphere sk y in the H a band.

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APPENDIX A NEAR-IR SUR VEY DET AIL OF OBSER V A TIONS The follo wing tw o tables contain the basic information for the nal selection of elds included for analysis in the present w ork, b ut not for all of the observ ations of e v ery single eld. V ariations of the quality of the data are notorious in tw o aspects of this database: the rst one is the a v erage seeing column of table A.1, which represents the median v alue of the stellar prole FWHM and which w as mostly af fected by weather conditions at the observ atory; the second aspect is the photometric scatter tab ulated in table A.2 as the median v alue of the standard de viations per magnitude of magnitude dif ferences with respect to 2MASS counterparts within the ranges 11 : 0 < J < 16 : 0, 11 : 0 < H < 15 : 0 and 11 : 0 < K < 14 : 0. These v alues are also af fected by weather conditions, b ut the y also contain information on intrinsic v ariations in the surv e y data quality due to, among others, v ariable focusing (a recurrent problem of the KPNO-2.1m telescope) and more recently (winter 2004) de gradation of the lens coating. Some of these aspects are also discussed in chapter 3 167

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168T able A–1. Summary of near -IR observ ations. FLAMINGOS KPNO-2.1m Field RA DEC Date Filter Exp Dithers T otal Exp. Airmass Seeing ID center J2000 (sec.) (sec.) (00) 1 6:31:53.36 4:59:16.8 2002 Jan 13 J 60 16 960 1.415 1.67 1 2002 Jan 13 H 60 16 960 1.311 1.72 1 2002 Jan 13 K 30 32 960 1.231 1.78 2 6:33:13.66 4:39:16.8 2003 Oct 19 J 60 16 960 1.148 1.61 2 2003 Oct 20 H 60 16 960 1.127 1.65 2 2003 Oct 20 K 20 50 1000 1.128 1.53 3 6:31:53.36 4:39:16.8 2003 Oct 19 J 60 15 900 1.131 1.63 3 2003 Oct 20 H 60 16 960 1.136 1.65 3 2003 Oct 20 K 20 47 940 1.150 1.55 4 6:34:05.45 4:24:54.2 2001 Dec 19 J 60 16 960 1.142 1.83 4 2001 Dec 17 H 60 14 840 1.256 2.16 4 2001 Dec 17 K 30 32 960 1.194 2.26 5 6:34:31.53 4:14:54.2 2001 Dec 19 J 60 16 960 1.132 1.88 5 2001 Dec 19 H 60 16 960 1.219 1.92 5 2001 Dec 19 K 30 31 930 1.287 1.63 6 6:33:13.59 3:59:16.8 2001 Dec 19 J 60 14 840 1.133 1.74 6 2001 Dec 19 H 60 15 900 1.184 1.77 6 2001 Dec 19 K 30 30 900 1.408 1.76 7 6:31:53.36 4:19:16.8 2001 Dec 19 J 60 15 900 1.136 1.78 7 2001 Dec 19 H 60 16 960 1.153 1.74

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169T able A–1—Continued Field RA DEC Date Filter Exp Dithers T otal Exp. Airmass Seeing ID center J2000 (sec.) (sec.) (00) 7 2001 Dec 19 K 30 29 870 1.537 1.91 8 6:35:51.76 3:54:54.2 2001 Dec 20 J 60 14 840 1.405 1.81 8 2001 Dec 20 H 60 16 960 1.318 1.81 8 2001 Dec 20 K 30 32 960 1.233 1.84 9 6:34:31.53 3:54:54.2 2001 Dec 23 J 60 16 960 1.390 1.92 9 2001 Dec 23 H 60 16 960 1.189 1.78 9 2001 Dec 23 K 30 30 900 1.146 1.93 10 6:33:13.59 4:59:16.8 2002 Jan 14 J 60 16 960 1.198 1.77 10 2002 Jan 14 H 60 16 960 1.161 1.94 10 2002 Jan 14 K 30 30 900 1.297 2.09 11 6:30:33.12 4:59:16.8 2002 Jan 14 J 60 16 960 1.126 2.08 11 2002 Jan 14 H 60 16 960 1.140 2.13 11 2002 Jan 14 K 30 32 960 1.204 2.08 12 6:29:12.88 4:59:16.8 2002 Jan 14 J 60 16 960 1.258 1.78 12 2002 Jan 14 H 60 15 900 1.142 1.96 12 2002 Jan 14 K 30 33 990 1.124 2.03 13 6:30:33.12 4:39:16.8 2002 Dec 20 J 60 16 900 1.451 2.08 13 2002 Dec 20 H 60 16 840 1.283 2.01 13 2002 Dec 20 K 30 30 900 1.221 1.91 14 6:33:13.59 5:19:16.8 2002 Dec 16 J 60 17 1020 1.131 1.76

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170T able A–1—Continued Field RA DEC Date Filter Exp Dithers T otal Exp. Airmass Seeing ID center J2000 (sec.) (sec.) (00) 14 2002 Dec 16 H 60 16 960 1.367 1.85 14 2002 Dec 16 K 30 32 960 1.175 2.08 15 6:31:53.36 5:19:16.8 2002 Dec 19 J 40 24 920 1.308 2.11 15 2002 Dec 19 H 60 16 960 1.219 2.13 15 2002 Dec 19 K 30 33 990 1.121 2.13 16 6:30:33.12 5:19:16.8 2002 Dec 19 J 60 15 900 1.409 1.87 16 2002 Dec 19 H 60 16 960 1.174 2.07 16 2002 Dec 19 K 30 29 870 1.141 1.93 17 6:33:13.59 4:19:16.8 2001 Dec 23 J 60 16 960 1.287 1.97 17 2001 Dec 23 H 60 15 900 1.231 1.85 17 2001 Dec 23 K 30 30 900 1.130 1.83 18 6:30:33.12 4:19:16.8 2003 Oct 22 J 60 16 960 1.134 1.61 18 2003 Oct 22 H 60 16 960 1.129 1.61 18 2003 Oct 22 K 20 60 1200 1.137 1.51 19 6:36:29.88 4:24:54.2 2003 Jan 30 J 15 60 900 1.296 1.97 19 2003 Jan 30 H 16 60 960 1.132 1.98 19 2003 Jan 30 K 35 28 980 1.155 1.89 20 6:35:51.76 4:14:54.2 2003 Jan 30 J 60 15 900 1.482 1.98 20 2003 Jan 30 H 60 16 960 1.136 1.84 20 2003 Jan 30 K 35 32 1120 1.270 1.88

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171T able A–1—Continued Field RA DEC Date Filter Exp Dithers T otal Exp. Airmass Seeing ID center J2000 (sec.) (sec.) (00) 21 6:35:51.76 3:34:54.2 2003 Jan 30 J 60 16 960 1.251 1.86 21 2003 Jan 30 H 60 16 960 1.161 1.91 21 2003 Jan 30 K 35 32 1120 1.191 2.07 22 6:34:33.83 4:39:16.8 2003 Jan 30 J 60 16 960 1.380 1.97 22 2003 Jan 30 H 60 16 960 1.126 1.94 22 2003 Jan 30 K 35 33 1155 1.199 1.89 G1 6:35:11.14 4:04:49.4 2004 No v 10 J 60 16 960 1.139 1.83 G1 2004 No v 10 H 60 16 960 1.135 1.71 G1 2003 No v 21 K 35 32 1120 1.141 1.58 G2 6:35:10.76 3:44:49.5 2003 No v 21 J 60 16 960 1.181 1.56 G2 2003 Dec 04 H 60 16 960 1.135 1.68 G2 2004 No v 09 K 35 28 980 1.255 1.71 G3 6:33:52.38 4:09:08.1 2004 No v 09 J 60 15 900 1.337 1.74 G3 2003 Dec 04 H 60 16 960 1.140 1.73 G3 2004 No v 09 K 35 30 1050 1.132 1.73 CF1 6:25:32.00 3:44:30.0 2002 Jan 05 J 60 16 960 1.329 1.88 CF1 2002 Jan 05 H 60 15 900 1.251 1.87 CF1 2002 Jan 05 K 60 19 1140 1.196 1.85 CF2 6:23:37.00 2:58:51.0 2004 Jan 27 J 60 16 960 1.149 1.92 CF2 2004 Jan 27 H 60 15 900 1.143 1.98 CF2 2004 Jan 27 K 35 32 1120 1.253 2.01

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172 T able A–2. Mean Photometric Scatter by Field Field s J s H s K 01 0.061 0.065 0.059 02 0.057 0.052 0.055 03 0.048 0.047 0.043 04 0.063 0.052 0.070 05 0.059 0.055 0.062 06 0.049 0.048 0.044 07 0.049 0.051 0.045 08 0.047 0.060 0.058 09 0.048 0.058 0.048 10 0.059 0.079 0.043 11 0.053 0.059 0.052 12 0.056 0.064 0.044 13 0.060 0.046 0.049 14 0.046 0.047 0.043 15 0.054 0.056 0.042 16 0.056 0.099 0.043 17 0.043 0.048 0.047 18 0.073 0.084 0.067 20 0.061 0.072 0.060 21 0.055 0.070 0.072 22 0.059 0.070 0.073 G1 0.071 0.086 0.082 G2 0.089 0.099 0.090 CF1 0.050 0.044 0.039 CF2 0.077 0.051 0.046 CF3 0.090 0.103 0.080

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APPENDIX B MILLIMETER SUR VEY DET AIL OF OBSER V A TIONS The follo wing table contains the total scan areas areas of the molecular emission maps constructed with the IRAM 30m data. The areas are e xpressed in square arcseconds and are or ganized by tracer as indicated in each column. T able B–1: IRAM Observ ations: Area Co v erage by T racer Map Area by tracer [sq. arcsec] Cluster CO(2-1) 13 CO(2-1) C 18 O(2-1) CS(2-1) HCO + (1-0) N 2 H + (1-0) PL01 38025 38025 38025 38025 38025 38025 PL02 45000 45000 47813 47025 PL03 50400 50400 19800 50400 50400 19800 PL04 47540 47540 12600 47540 47540 12600 PL05 40500 40500 11813 40500 40500 11813 PL06 43594 43594 34650 43594 43594 34650 PL07 41850 41850 22275 41850 41850 22275 RLE08 41625 49218 47194 49218 49218 25988 1 1 emission detected only in southern latitude clump (RLE08B) 173

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APPENDIX C MILLIMETER SUR VEY DET AIL MUL TIP ANEL MAPS This appendix contains a detailed discussion for each one of the indi vidual multipanel molecular emission maps of Figures 5–1 to 5–8 Reference to nomenclature of IRAS point sources and Bell Labs clumps follo w the studies of PL97 and WBS95. Cluster PL01: Associated with clump 11 from WBS95 and IRAS 06291+0421; WBS 13CO maps sho w that the clump is separated from the main body of the cloud in the 15 to 120 Km/s range. The C 18 O(2-1) emission w as weak b ut the other tracers sho wed strong, compact emission well coincident with the cluster HCO + (1-0) and CS(2-1) emission peaks present a 0.5 : 0 of fset (0.23 pc) which is signicant gi v en the resolution of the maps. Cluster PL02: Associated with clump 18 from WBS and IRAS 06306+0437. The morphology of the cluster itself is interesting because of its “T” shape structure, with the top of the T suggesting inuence of the ionization front (molecular cloud ridge ) (see section 2.3.1 Contours appear to compress dramatically to w ards the SE re gion, roughly in the direction of the neb ula. The surf ace density of the cluster is actually lo wer than the a v erage. The N 2 H + (1-0) and the C 18 O(2-1)emission to w ards this cluster were e xtremely weak, to the point that we were not able to construct maps for these tw o tracers. The HCO + (1-0) emission is almost twice as strong as CS(2-1) Cluster PL03: second to furthest from the neb ula, associated with clump 7 from WBS and IRAS 06308+0402, the clump is part of a se gment separated from the main cloud body in the range 15 to 120 km/s. The HCO + (1-0) and CS(2-1) intensity peaks present a notorious 1.8 0 of fset (0.8pc), with HCO + (1-0) a possible signature of molecular dif ferentiation. HCO + (1-0) peaks near 174

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175 the south part of the cluster while CS(2-1) and N 2 H + (1-0) peak at a group red sources in the SW edge. Cluster PL04: corresponds to clump 1 from WBS and IRAS 06314+0427. The core contains about 40 sources which are embedded in a red neb ulosity b ut the cluster is lar gely e xtended be yond this core and mix es with the population of cluster PL05. The stars-gas conguration in PL05 is v ery similar to PL04. CO(2-1) and 13 CO(2-1) emission 4 is distrib uted in a lar ge plateau, with a wide peak at the red core. The “dense” tracers sho w the opposite situation, with lo w intensity at the peak CO(2-1) and 13 CO(2-1) and stronger emission at a second stellar density peak located NW N 2 H + (1-0) and C 18 O(2-1) also sho wn the ”strongest” intensity at this NW zone, whereas it w as v ery f aint e v erywhere else (hence the v ery small maps obtained). In summary the most obscured zone of the cluster is practically v oid of dense gas, while the zone that appears clean of neb ulosity (dusty gas) is where dense tracers tend to accumulate. Cluster PL05: The situation is similar to that of cluster PL04. Dense gas emission peaks locate near a v ery embedded group of stars SE of the main stellar density center This situation applies for all of the six tracers, e xcept for 13 CO(2-1) which hints a ne w condensation to w ard the NE where cluster PL04 is located. There is an moderate of fset (0.14 pc) between the centers of emission of CS(2-1) and HCO + (1-0). Cluster PL06: Possibly younger than the rest: hosts IRAS 06319+0415, the most luminous in the re gion, which is coincident with AFGL961, B PMS binary with an associated CO outo w and possible self-absorption. All tracers and sho w compact, well localized emission that forms an elongated clump which e xtends in the NW -SE direction. The peak of emission in all tracers is coincident with the source AFGL961W which is kno wn to ha v e a bo w shock. This cluster presents no of fset between peaks of emission, which denies molecular dif ferentiation.

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176 If the cluster is, younger than those with tracer peak of fsets, then the ef fect is clearly e v oluti v e. Cluster PL07: Associated with clump 7 from WBS and IRAS 06329+0421. This cluster is the least luminous, the least dense and furthest from the Rosette Neb ula in the sample. In terms of number of members and size, it is v ery similar to PL06 The morphology of all tracers is compact and well localized, with the main emission running from NE to SW in the form of a compact, well dened elongated clump. The peak emissions of CS(2-1) HCO + (1-0) and N 2 H + (1-0) are separated from the cluster density center by about 1 arcmin in the a v erage (0.46 pc) while the of fsets of C 18 O(2-1) and 13 CO(2-1) are almost ne gligible. Cluster RLE08A: Associated with the clump 17 from WBS95 and ”f aint” IRAS source 06314+0421, this is the most disperse cluster of the set, with an e xtension of almost 1.6 parsecs in a v erage radius from the center of maximum density The emission spectra in the area sho ws tw o well separated components, with maxima at around 11.0 (clump 8a1) and 16.0 km/s (clump 8a2) respecti v ely The clump 8a1 feature has an ample rms dispersion (more than twice than its companion, in e v ery tracer), resulting in a loose clump. The feature associated with clump 8a2 is a v ery narro w high v elocity component, and the resultant clump is compact and possible closer to equipartition. The CO(2-1) and 13 CO(2-1) emission has a high le v el of rms noise, and the clumps are dif cult to dif ferentiate. It is only in CS(2-1) and HCO + (1-0) where we can distinguish the tw o separate cores. The N 2 H + (1-0) emission w as v ery weak in the whole area e xcept at the re gion 8a2, where a v ery small, compact clump w as found. The C 18 O(2-1) emission map w as too noisy and weak lik e in PL02 and we could not construct a reasonable map for that tracer

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BIOGRAPHICAL SKETCH Carlos Rom an-Z u niga (1972) w as born in Me xico City According to the people who kno w him, he is a natural astronomer mainly because of his space y beha vior Some of his relati v es say that he ne v er has been able to set his feet on the ground completely and this characteristic al w ays mortied his f amily specially his parents. According to a local le gend, Carlos started lo ving astronomy after w atching Carl Sagan' s Cosmos series on TV during his childhood, b ut some other testimonies point to a physics course in high school taught by a v ery enthusiastic professor H ector P adilla. Apparently one day he e xplained the La ws of K epler and Carlos w as the only one among a group of ya wning teenagers who w as drooling with scientic allure. After that life-changing e xperience, Carlos decided to study physics at the Uni v er sidad Nacional Aut onoma de M exico (UN AM). F or his good luck, this w as a physics program where he could tak e electi v e astrophysics courses with the professors of the Instituto de Astronom a, one of the lar gest astronomy research centers in Latin America. At UN AM, Carlos not only conrmed his lo v e for the stars and the la ws that go v ern the w orld, b ut he also de v eloped a “beatnik” taste for science, which manifested later in a casual –b ut ne v er supercial– passion for understanding his surroundings (dened as a sphere of radius 1 Gpc). In 1995, Carlos went studying abroad, and he enjo yed six months of perfect weather at the Uni v ersity of California in Santa Cruz. There he met Michael Bolte, an e xpert in the study of glob ular clusters, who of fered him a small project. This collaboration resulted in a v ery successful Licenciatura thesis, which allo wed him to obtain a Physicist de gree and opened the doors of graduate school. 183

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184 Needless to say after this happ y episode Carlos' lo v e for the stars only gre w lar ger to the dismay of his f amily who probably w ould prefer him to pursue a more do wn to Earth career The reasons of course, had nothing to do with mone y or goals, b ut with the hope that he could remember to tie his shoes from time to time. Ignoring (or maybe for getting) these w orries, Carlos decided to enter a graduate program in astronomy After a good year and a half of hard course w ork in the master' s program at the Instituto de Astronom a, UN AM, he nally took the decision of mo ving to Florida to get closer to the big telescopes. In Florida he met Elizabeth Lada, a f amous astrophysicist who told him ho w f antastic w as to observ e clusters of stars when the y were still young, and she in vited him to join the eld of star formation. Little he kne w at that time he w ould spent o v er 7 long b ut intense years on this rst di v e in the ocean of professional astrophysics. F ortunately for Carlos (and for the peace of mind of his parents), his belo v ed wife F abiola joined him in this trip, and stayed at his side all that time. W ell, almost, because last semester she mo v ed to Amherst, Massachussetts, to start her o wn PhD in Anthropology and also to put some pressure on Carlos to nish his thesis. Carlos will mo v e to Massachussetts too after defending his dissertation, and he will w ork for a little while at the Center for Astrophysics in Harv ard. After that he plans to sit do wn on a bench looking for Astronomer W anted adds in the ne wspapers, al w ays with the dream of k eep studying the stuf f be yond the clouds for as long as possible. Being a professional astronomer is a good and big step for Carlos, and his parents and wife are really proud b ut, according to testimonials, the y mostly hope the de gree can help him to remember to turn of f the sto v e or to iron his shirts re gardless that his mind is se v eral hundred parsecs a w ay


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Title: Near-Infrared Study of the Star-Forming Properties of the Rosette Complex
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Copyright Date: 2008

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Title: Near-Infrared Study of the Star-Forming Properties of the Rosette Complex
Physical Description: Mixed Material
Copyright Date: 2008

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NEAR INFRARED STUDY OF THE STAR-FORMING
PROPERTIES OF THE ROSETTE COMPLEX















By

CARLOS G. ROMAN-ZUNIGA


A DISSERTATION PRESENTED TO THE GRADUATE SCHOOL
OF THE UNIVERSITY OF FLORIDA IN PARTIAL FULFILLMENT
OF THE REQUIREMENTS FOR THE DEGREE OF
DOCTOR OF PHILOSOPHY

UNIVERSITY OF FLORIDA


2006


































Copyright 2006

by

Carlos G. Romin-Zifiiga


























This work is dedicated to the memory of:


Richard J. Elston (1961-2004)


and


Leonel Hernandez (1972-2001)















ACKNOWLEDGMENTS

This is the second time in my life I have to write an acknowledgment section for

a thesis. Just as then, I will make my best effort to to avoid omitting important names.

However, if I do, let me say -as some sort of disclaimer- that it was not on purpose. My

memory is very selective and it tends to retain too much movie trivia, old jingles and bad

jokes, while randomly erasing names, telephone numbers and birthdays.

But seriously:

I want to acknowledge in a general way the Department of Astronomy at the Uni-

versity of Florida accepting me as a student and for giving me a Teaching Assistantship

during my first two and a half years in graduate school. Also, for giving me an office to

study, access to a computer, free photocopying services and a marvelous work environ-

ment.

I want to acknowledge CONACYT-Mexico for a fellowship that sponsored a major

fraction my doctoral studies at the University of Florida. During 4 years CONACYT gave

me a living stipend, paid my tuition in full and paid for a significant part of my health

insurance costs. My most sincere gratitude to this excellent program.

It is also important to mention that FLAMINGOS, the instrument we used to collect

most of our data was designed and constructed by the IR instrumentation group (PI: R.

Elston) at the University of Florida, Department of Astronomy, with support from NSF

grant AST97-31180 and Kitt Peak National Observatory.

Next, I want to thank first my supervisor, Dr. Elizabeth Lada. Her efforts to direct

my graduate career have been enormous to say the least, and I feel that it is only fair

to say that without her, none of this would have been possible. When I asked her for a

summer job at the end of my first year at UF, she told me about this 'nice little project










on the Rosette Molecular Cloud, and how she would be happy to have me in charge

of it. Well, I took the summer project, the project grew to a thesis, and I fell in love

with it completely. Six years later, I am still working on it, pondering about its many

consequences, and hoping to keep working on it for a while. You see, I already traveled

to the Rosette in my dreams, and I got to know the place well. Yes, there were many

challenges along the way, both technical and personal, but I cannot but admire Elizabeth

for always being so enthusiast -and for being patient with my own flashes of over-

enthusiasm-, for never allowing me to give up no matter how discouraging were the

problems, and mostly for being there, sometimes as a supervisor, sometimes as some sort

of basketball couch, sometimes as a psychologist, but most of the time as a friend. On top

of that, Elizabeth generously used part of her grant money to take care of my salary and

tuition requirements during my last year and my first two summers.

Dr. Richard Elston constructed the instrument FLAMINGOS it with the skills and

patience of a clockmaker, and then he made sure we cared about it the same way he did.

I remember those first two years of the survey, where the goals were still confusing and

foggy, our pipelines too buggy and the piles of data overwhelming: Richard always had a

simple and clear way to solve any problem. Then he got ill and had to leave us, but I am

glad I had enough time to learn from him a very important lesson, one about keeping up

the courage intact even in the most difficult of circumstances.

On another comer of the FLAMINGOS project is the effort and amazing stamina of

Dr. Nick Raines, who not only knows every single cable and bolt of the instrument, but

also every inch of the telescope facilities and every fine dinning corner of Tucson.

I want to thank Dr. Jonathan Williams, from the Institute for Astronomy at the

University of Hawaii, and former professor of Astronomy at UF. He took the necessary

time to direct one entire chapter of this thesis, showing me the world of radioastronomy

all the way to the big leagues at FCRAO and IRAM. And because he is an expert on the

Rosette, there were many things I learned straight from his papers.










I want to thank the faculty of the Astronomy Department for their patience, their

help, their classes, their advice and their restless effort to integrate what is on its way

to be one of the best Astronomy programs in the world. I want to mention Dr. Stanley

Dermott, chairman of the department, who gave me trust and support even in difficult

academic moments, Dr. Francisco Reyes for being such a great supervisor as coordinator

of the Astronomy Laboratory, and a great example to follow; Dr. Rafael Guzman for

being always a supportive friend and informal advisor, and Dr. Ata Sarajedini for always

being attentive to my career as Graduate Coordinator.

To the secretary staff of the Astronomy Department I wish I could dedicate a

chapter of my thesis. There are no words to describe their help, specially with all the

extra paperwork that my condition of foreign student implies. I want to specially thank

Catherine Cassidy for dedicating so much time to remind me of academic deadlines,

always with that extra wit and animosity. Debra and Deborah for all that help with grants,

assistanships and travel, Glenda for helping me with my employee documentation, and

also Tracey and Ann, who are in different departments now, but always were beautiful

and patient and showed me how to fight the bureocratic monsters.

Joanna Levine was my officemate, my project colleague, and a wonderful friend.

Joanna is the kind of person that you can talk to about almost anything, and believe me, in

5 years of sharing an office with someone you get to cover a lot of conversation material.

Now I really hope time will keep us being friends and collaborators, even if she becomes

the first one in the world to combine astronomy and dance into a single, indivisible

discipline and become insanely famous.

Thank you to Dr. August Muench, also ex-officemate, for his brilliant lessons on

computing, scientific passion and insomnia. Many ideas in this thesis became clear

after commenting them with him. Other "ideas", for the same reason, went safely to the

garbage pail before being incorrectly stated. Now we will be officemates again, which is a

good opportunity to learn to follow very high standards like he always does.










I also want to thank Dr. Charles Lada, from the Harvard Smithsonian Center for

Astrophysics for many useful comments and encouraging words about this project, some

of them crucial for its completion. And nowadays for giving me an opportunity to work

with him. Astronomy does not get much better than this.

I want to thank Eric McKenzie for his patience coping with the many versions of

the pipeline and to enjoy observing and reducing star formation data just as he enjoys

observing and reducing those long integration plates of galaxy clusters. And thanks for

his anecdotes about walking across the United States, reading a thousand books and the

nutritional virtues of peanut butter.

I want to thank Dr. Anthony Gonzalez (nowadays a professor at UF), Dr. Matthew

Horrobin, Dr. Andrea Stolte and Dr. Aaron Steinhauer, all former postdoctoral fellows,

for taking our graduate student mess and helping to convert it into a refined scientific

effort. At the mountain, they mastered the many twiggles of FLAMINGOS and the

KPNO telescopes. Back at the office, they were always available for questions and

comments, and I cannot remember one single time when the answers were not given

to us with a big side of smiling: Anthony always was inquisitive enough to discover

obscure bugs in the pipelines, but also patient enough to wait for us to correct them.

With Matthew I learned tons too, and shared with him the experience of a data quality

assessment trip to Boston, complete with bar hopping and all. From Aaron I will always

admire his neat style and enthusiasm, everything to him was an opportunity to learn,

which is fantastic. Oh, and his Simpsons collection kept us awake many times at the

telescope! Finally, Andrea was to put it in simple words, the person who saved my

project. She was patient enough to understand the problems we were dealing with and

then even more patient to fix them. She even accepted my crazy idea of making a catalog

software application in Supermongo and then refined the effort!. That is something to

admire. Really.










I want to thank Noah Rashkind and Chris Foltz for being the guinea pig users of

LongLegs, with all that amazing vitality and enthusiasm. They were also great trip pals in

the Boston trip, which was a lot of fun.

Many thanks to Bruno Ferreira and Jorge Gallego. Bruno for instance, made possible

a crucial gear of my analysis, digesting successfully a cryptic little paper I reccomend him

to read. But that is nothing compared with the many many great moments we shared as

part of the Star Formation Crib, and many friendly gatherings. Jorge, the other creature

lurking in the cave 319 always kept me afloat with his cheerful style, along with many

discussions about movies, spanish rock bands and the meaning of life. Nowadays these

guys are so important in the Gainesville community (specially Maestro Yogi Bruno) that I

wonder how the city will cope with after they leave. Thanks also for pizza, for the Indian

Air sessions, for FILMINGOS and for many other adventures.

I want to thank the many UF graduate students that I met along the years. Names

like Lauren, Elisha, Christos, Sue, Barbara, Jim, James, Scott, Rob and TJ might sound

ancient to some people now, but they were the same as the rest of us not a long ago. In a

second layer are Kelly, David, Debbie, Doug, Pimol, Paty, Veera, Derrick, Bill, Catherine

and Manuel among those who left. Craig, Gator extraordinaire is a separate case. And

then all of those who keep the joy going on: Eric, Ana, Ileana, Naibi, Cynthia (thanks

for the hospitality!), Suvrath, Margaret (the coffee hour angel), Michelle, David, Aaron,

Mike, Ashley, Valerie, Paola, Miguel, Sung, Justin, Lauren, Curtis, Scott, Audra, Alison,

Justin, Leah, Julian and Andrew.

To the Mexicans in Gainesville Student Association, for being my family all these

years. Without their support, adapting to a whole new country would have been near

to impossible. I have a special mention for Julio Castro for being my lunch pal, joke

sidekick, movie critic partner and friend all these years. I also want to thank Eugenio y

Milena (and her parents Sofia and Fernando), Velia y Luis, Diego y Erica, Jorge, Rocio,

Juan, Hussein, Alicia, Maria Jose y Leo, Horacio y Maru, Arturo y Rosa Isela, Antonio










y Roxanna, Sebastian y Paula, Nicasio y Miriam. I owe an apology for not putting the

names of their children -I am already several pages above expected-. Too many moments

we shared: meetings, barbecues, carnivals, September 15th parties, birthdays, you name

it, but the best thing are the memories.

To my family, for ALWAYS believing in me and being close to me despite distance

and the time. My parents, Hector and Rosario have always been my greatest motivation,

and everything I am now I owe to them and their efforts and many sacrifices. To my

brother and sister, Esteban and Daniela, I want to say that we are all part of the same love,

and I am proud of you every moment. I love you all so much, and I need to be with you

again so much... I only pray for that moment to come soon.

To beloved Sheikha Amina Teslima and all the members of the Al-Jerrahi commu-

nity for keeping the very essential light of my heart always lit on. And for making time

and distance invisible. Alhamdulilah.

And Finally, the person who I decided to share the road of life with. Fabiola, can you

believe that I am writing this for the second time around? And how many words do I need

now to explain what I feel, if there are infinite reasons for being thankful and happy? No,

I cannot start nor finish because you are my beggining and my end. I love you with every

part of me.
















TABLE OF CONTENTS
page

ACKNOWLEDGMENTS ............................ iv

LIST OF TABLES ...... ................. ........ xiii

LIST OF FIGURES ........... ...... ........ ........ xiv

KEY TO ABBREVIATIONS ............................. xviii

KEY TO SYMBOLS .. ............................ xix

ABSTRACT. ..................... ................. xx

CHAPTER

1 INTRODUCTION ................... ......... 1

1.1 A Global Picture of Star Formation .......... ....... ... 1
1.2 Motivations for the Study of the Rosette Complex ............ 4

2 THE ROSETTE COMPLEX IN MONOCEROS ........ ......... 7

2.1 Historical Perspective ........................... 7
2.2 The Rosette Nebula and the Young Cluster NGC 2244 ......... 9
2.2.1 The Rosette Nebula ......... ............... 9
2.2.2 NGC 2244 ............................. 10
2.2.3 Spectroscopic Studies .......... ....... ..... 12
2.2.4 Near Infrared Studies ......... ............. 13
2.2.5 X-ray Studies ................... ...... 14
2.3 The Rosette Molecular Cloud: Structure ....... ......... 14
2.3.1 CO studies ............................. 14
2.3.2 Interaction with the Rosette nebula ......... ....... 18
2.4 The Rosette Molecular Cloud: Embedded Populations . ... 23
2.4.1 Dominance of Cluster Formation in the Rosette Complex . 23
2.4.2 The Hypothesis of Sequential Star Formation . ... 25

3 A NEAR-IR SURVEY OF THE ROSETTE COMPLEX: OBSERVATIONS 27

3.1 The FLAMINGOS GMC Survey ....... . .... 27
3.2 Data Reduction ........... . . . ......29
3.2.1 The Data Reduction Pipeline: LongLegs . . ... 29
3.2.2 The Photometry and Astrometry Pipeline: PinkPack ....... .31










3.3 Completeness of Sample .................. ..... 32
3.4 Positional Correction of Photometry . . ...... 35
3.5 Quality and Uniformity of the Survey .... . . . 40
3.6 Construction of Final Catalog ...... . . . .... 44
3.6.1 Intrinsic quality: 2MASS Addendum . . ... 44
3.6.2 Survey Area Merging ........ . . .... 44
3.7 Intrinsic Detection Constraints .............. . 45

4 NEAR-IR SURVEY: ANALYSIS AND RESULTS . . . 48

4.1 Introduction .............. . . ..... 48
4.2 Analysis .......................... .... ........ 50
4.2.1 The Nearest Neighbor Method . . 50
4.2.2 Detection of Embedded Populations . . . 52
4.2.3 Infrared Excess Stars .................. ..... .. 53
4.2.4 Magnitude Depth Restriction for IRX stars . . .... 55
4.2.5 Nearest Neighbor Analysis for Infrared Excess Stars ....... 58
4.2.6 Identification of Clusters .............. ... .. 60
4.2.7 Properties of Clusters ................ .... .. 63
4.3 The Fraction of Stars in Clusters . . . .... 69
4.3.1 Distribution of Sources with Respect to the Rosette Nebula 75
4.3.2 A Case for a Distributed Population? . . . 79
4.4 Discussion and Future Work ................ .... .. 82

5 OBSERVATIONS OF CLUSTER DENSE GAS ENVELOPES ......... 87

5.1 Introduction ............... . . . 87
5.2 Observations and Data Reduction ..... . .... 89
5.3 Analysis and Results.. . . . ..103
5.3.1 Presentation of the Data ....... . . ... 103
5.3.2 Local Extinction ......... . . . ... 104
5.3.3 Calculation of Physical Parameters . . . .. 105
5.3.4 Tracer Abundances . . . ..106
5.3.5 Gas Mass ....... ..................... 108
5.3.6 Clump Sizes .................. ........ 109
5.3.7 Line W idths ................... . .. 110
5.3.8 Clump Masses .......... . . ... 110
5.3.9 Gas Dynamics .......... . . ... 114
5.4 The Embedded Stellar Population .... . . .... 116
5.4.1 Star Forming Efficiencies . . . 116
5.4.2 The Gas Stars Connection ..... . . .. 117
5.5 Chemical Differentiation in Cluster Envelopes . . ... 119
5.6 Summary and Discussion .................. ...... ..121










6 GLOBAL ASPECTS .......... ....................... 125

6.1 A Near Infrared Extinction Map for the RMC . . ... 125
6.1.1 Motivations ..... . . . . . 125
6.1.2 Dust Extinction from Near-Infrared Colors: NICE and NICER 126
6.1.3 The Extinction Map . ....... . . 130
6.2 Individual Extinction Cores .............. ..... 134
6.2.1 Identification and Estimation of Properties . . .... 134
6.2.2 Core Sizes ............ . . .... 143
6.2.3 Core M asses ................... . .. 145
6.2.4 The Embedded Cluster Mass Function . . ... 148
6.2.5 Star Formation Efficencies .... . . .. 153
6.3 Summary and Discussion .................. ...... ..157

7 CONCLUSIONS AND FUTURE WORK ..... . . .. 161

7.1 Distribution of Young Stellar Populations in the Rosette Complex .. 161
7.2 The Local Environments of Young Clusters . . .. 162
7.3 Extinction in the Rosette Complex and Global Results . ... 163
7.4 Future Work .................. ............. 164

APPENDIX

A NEAR-IR SURVEY. DETAIL OF OBSERVATIONS . . ... 167

B MILLIMETER SURVEY. DETAIL OF OBSERVATIONS . . ... 173

C MILLIMETER SURVEY. DETAIL MULTIPANEL MAPS ..... ... .. 174

REFERENCES ................... . . .... 177

BIOGRAPHICAL SKETCH .......................183
















LIST OF TABLES


Table

2-1

4-1

5-1

5-2

5-3

5-4

6-1

6-2

A-1

A-1

A-1

A-1

A-2

B-1


page



. . 64
. . 64


. . 102

. . 106

. . 116

. . 139

. . 157

. . 168

. . 169

. . 170

. . 171

. . 172

. . 173


Distance Estimates to the Rosette (NGC 2244) . ...

Young Clusters Rosette Complex . ..........

Relevant Properties of Rosette Clusters ............

Clump properties for Rosette Clusters . .

Molecular Line Parameters . .............

Star Formation Efficiencies (13CO(2-1)) . ......

Extinction Cores in the Rosette Complex ...........

Star Formation Efficiencies (Av cores) . .......

Summary of near-IR observations. FLAMINGOS KPNO-2.1m

Summary of near-IR observations. FLAMINGOS KPNO-2.1m

Summary of near-IR observations. FLAMINGOS KPNO-2.1m

Summary of near-IR observations. FLAMINGOS KPNO-2.1m

Mean Photometric Scatter by Field . .........

IRAM Observations: Area Coverage by Tracer . ..
















LIST OF FIGURES
Figure page

2-1 A photograph of the Rosette Nebula .................. .... 7

2-2 Location of the Rosette Cloud in the Monoceros Complex . . 8

2-3 A Ha vs. V-I diagram for NGC 2244 .................. ..13

2-4 A CO map of the Rosette Molecular Cloud ... . . ..... 16

2-5 Molecular Clumps in the Rosette Cloud .................. ..17

2-6 IRAS 12pm and 1400 Mhz map overlay ..... . . 20

2-7 A 0.5-2 keV Chandra image of the Rosette Complex . . ... 22

2-8 Location of the Phelps & Lada clusters ................. ..24

3-1 UF/NOAO Rosette Complex Survey Map ................. ..30

3-2 Completeness Limits by Filter .................. ....... 33

3-3 Completeness by type of field ..... ........ ...... 34

3-4 Photometric Quality Areas .................. .. ...... 37

3-5 Photometric Correction: Radial (J,H) .................. ..38

3-6 Photometric Correction: Radial (K) ................... . 39

3-7 Photometric Correction: Colors .................. ....... 40

3-8 Photometric Correction: Color and Magnitude Diagrams . ... 41

3-9 Photometric Correction: Photometric Scatter ..... . . . 42

3-10 Distribution of Photometric Uncertainties by Filer . . .... 43

3-11 Detectability of an Embedded Populations ................. ..46

4-1 Areas of the Color-Color Diagram .................. 54

4-2 Contour level J- H vs. H- K Diagram for All stars in the Survey . 56

4-3 Contour level J- H vs. H- K Diagrams Divided by Brightness ...... 57

4-4 Nearest Neighbor Distributions for Bright IRX Stars . . .. 59










4-5 Location of IRX Stars with Brightness K < 15.75 mag .....

4-6 Identification of Clusters in the Rosette Complex ..............

4-7 Distribution of Cluster Core and Total Radii .. ..............

4-8 Analysis Plots for Cluster PL01 ........................


Analysis Plots

Analysis Plots

Analysis Plots

Analysis Plots

Analysis Plots

Analysis Plots

Analysis Plots

Analysis Plots

Analysis Plots

Analysis Plots


for Cluster PL02 ........................

for Cluster PL03 ........................

for Cluster PL04 ........................

for Cluster PL05 ......................

for Cluster PL06 ........................

for Cluster PL07 ........................

for Cluster RLE08 .. ...................

for Cluster RLE09 .. ...................

for Cluster RLE10 .......................

for Cluster NGC 2237 .....................


4-19 Analysis Plots for Cluster NGC 2244 .....................

4-20 Distribution of IRX stars as a function of distance to the Rosette Nebula

4-21 Cumulative Counts of IRX sources in Field 09 ...............

4-22 Images of Distributed Formation in Field 09 of the Survey ..........

5-1 Molecular Emission Maps: Cluster PL01 ...................

5-2 Molecular Emission Maps: Cluster PL02 ...................

5-3 Molecular Emission Maps: Cluster PL03 ...................

5-4 Molecular Emission Maps: Cluster PL04 .. ..............

5-5 Molecular Emission Maps: Cluster PL05 ...................

5-6 Molecular Emission Maps: Cluster PL06 ...................

5-7 Molecular Emission Maps: Cluster PL07 ...................

5-8 Molecular Emission Maps: Cluster RLE08A .................

5-9 Extinction in 13CO(2-1) Map Areas (1) ....................










5-10 Extinction in 13CO(2-1)Map Areas (2) .

5-11 Abundance Ratios by Tracer . .

5-12 Distribution of Clump Sizes by Tracer

5-13 Distribution of Line Widths by Tracer

5-14 Distribution of Virial Mass by Tracer .

5-15 Distribution of LTE Mass by Tracer ...

5-16 Distribution of Virial to LTE Mass ratios

5-17 Comparison to Curves of Binding Pressur

5-18 Distribution of Velocity Gradients by Trac

5-19 Distribution of Star Formation Efficiencie

5-20 Correlation Between Cluster Sizes and Er

5-21 Overlap of HCO+(1-0) and CS(2-1) Emis

6-1 Near-Infrared Extinction Map of the Rose

6-2 13CO emission map of the Rosette Compl

6-3 Contour Extinction Map with Cluster Pos

6-4 Correlation between 13CO and Av .

6-5 Distribution of individual Av values betw

6-6 Extinction maps for Individual Cores (1)

6-7 Extinction maps for Individual Cores (2)

6-8 Extinction maps for Individual Cores (3)

6-9 Extinction maps for Individual Cores (4)

6-10 Extinction Core Profiles (1) . .

6-11 Extinction Core Profiles (2) . .

6-12 Extinction Core Profiles (3) . .

6-13 Extinction Core Profiles (4) . .

6-14 Distribution of Core Radii . .

6-15 Comparison of Cluster and Core Radii .


. . . 100

.. . 107

. . 109

....... . ... 110

..... . . .. 111

........ . . .. 112

...... . . . 113

e ..... . . . 114

cer . . . 115

s ..... . . . 117

mission Offsets . . 118

sion for Cluster PL03 . 119

:tte Complex . . 128

ex . . . 129

itions .. . . 131

. . .. . . 132

een0 and30mag . ... 133

..... . . . 136

..... . . . 137

..... . . . 138

..... . . . 139

. .. . 140

. . . . . 14 1

. . . . . 142

. . . . . 143

. . . . . 144

....... . . 145










-16 Cluster Radii and IRXF vs. Extinction . .

-17 Mean Extinction vs. Core Radii . ....

-18 Distribution of Extinction Core Masses . .

-19 Distribution of Core Mass Compared to Clusters .

20 The Embedded Cluster Mass Distribution Function

21 Rosette Clusters in the ECMDF . ....

22 Extinction Cores SFE vs. Cluster Radii . .

23 SFE as a Function of Distance to NGC 2244 .

24 Schematic Map of the Rosette Complex . .


. . . 146

. . . 147

. . . 148

. . . 149

. . . . 150

. . . 15 1

. . . 154

. . . 155

. . . .. 156


xvii
















KEY TO ABBREVIATIONS


2MASS

ECMDF

FCRAO

FITS

FLAMINGOS
trometer

FOV

GMCs

HBL

IMF

IRAF

IRX

ISM

NNM

NOAO

OBAs

PMS

PSF

ROSAT

SFE

SQIID

SSF

ZPT


Two Micron All Sky Survey

Embedded Cluster Mass Distribution Function

Five College Radio Astronomy Observatory

Flexible Image Transport System

Florida Multi-object Imaging Near-IR Grism Observational Spec-


Field of View

Giant Molecular Clouds

Hydrogen Burning Limit

Initial Mass Function

Image Reduction and Analysis Facility

Infrared Excess

Interstellar Medium

Nearest Neighbor Method

National Optical Astronomy Observatory

OB Associations

Pre-main sequence star

Point Spread Function fitting method

Rosetta X-Ray Satellite

Star Formation Efficiency

Simultaneous Quad Infrared Imaging Device

Sequential Star Formation

Photometric calibration zero point


xviii















KEY TO SYMBOLS


avir Ratio of Virial to LTE clump mass

Av Visual Extinction

E(B-V) Visual Color Excess

Ha Ha emission

HII Ionized Hydrogen

J, H, K Near Infrared Bands at 1.2, 1.6 and 2.2 pm

Mcc Mass of star forming extinction core

Mclus Cluster Mass

Memb Mass of Embedded Stars

MLTE Clump LTE mass

Msc Mass of starless extinction core

M. Units of Solar Mass

Mvir Clump virial mass

Nemb Number of Embedded Stars

Rcc Radius of Extinction Core

Rcore Cluster Core Radius

Req Cluster Equivalent Radius

Rv Visual Extinction to Excess Ratio















Abstract of Dissertation Presented to the Graduate School
of the University of Flr 'iida in Partial Fulfillment of the
Requirements for the Degree of Doctor of Philosophy

NEAR INFRARED STUDY OF THE STAR-FORMING
PROPERTIES OF THE ROSETTE COMPLEX

By

Carlos G. Rom6n-Ztfiiga

May 2006

Chair: Elizabeth A. Lada
Major Department: Astronomy

The Rosette Complex is one of the most important astrophysical laboratories for the

study of star formation. In this region we can study the interaction of an expanding HII

region -impulsed by the stellar winds from the large OB association NGC 2244- with

a large remnant molecular cloud, which is known to host seven embedded clusters. As

part of a large observational program to study the nature of young stellar populations in

giant molecular clouds, we made a complete near-infrared imaging survey of the Rosette

Complex using the detector FLAMINGOS. This survey is deep enough to detect stars

near the brown dwarf limit, improving considerably over available databases.

However, given the location of the Rosette Complex at a large distance from the

Sun and at a latitude close to the galactic disk, the contamination of the survey data by

field populations is high. In order to facilitate the detection of young populations, we

combined a selection of cloud members by means of their infrared excess emission with

a technique to detect star clusters using distances to nearest neighbors. This way we were

able to confirm the seven clusters previously identified, and to discover four new clusters.










For every stellar cluster we determined for the first time their approximate extensions

and number of members. We found that the fraction of stars in clusters in the Rosette

Complex is close to 87%, which is similar to other clouds like Orion. However, the

formation of clusters in the Rosette seems to be heavily influenced by the interaction

with the expanding nebula, as evidenced by the fact that the core of the molecular cloud,

coincident with the shock front of the expanding nebula contains 50% of the total cluster

population. The clusters in the core are also more extended and more populated.

Our study was complemented with a high resolution millimeter wavelength radio

survey of the dense gas emission around the 8 most prominent clusters in the sample.

We confirmed that all of the clusters observed are still embedded in what appear to be

very compact parental clump remnants, but in many cases these gaseous envelopes are

possibly becoming gravitationally unbound, due to the partial emergence of the young

cluster stars. The dense gas maps show features characteristic of the interaction of

clusters their local environment, particularly significant offsets of tracer emission peaks,

possibly due to chemical differentiation effects.

Our near-infrared observations also allowed us to construct an extinction map for

the fields observed. The map shows an good agreement with 13CO emission radio maps,

and allowed us to identify the main molecular cores in the complex. Using the mass

of stars in the clusters and the mass of the emission cores we calculated star formation

efficiencies, which resulted to be significantly larger at the central core of the cloud. Also,

extinction appears to be inversely proportional to the size of the clusters, but directly

proportional to the fraction of IRX sources, which is suggestive of evolutive effects and

a rapid dispersion of the gas after clusters are formed. The cluster emergence time scales

could be similar and even shorter than the T Tauri phase of the stars.















CHAPTER 1
INTRODUCTION

1.1 A Global Picture of Star Formation

Star Formation is one of the main puzzles in present day Astrophysics. Along

the years, it has been possible to construct a relatively detailed picture of the physics

involved in the formation of individual stars (Shu et al., 1987), but the problem of how

to extrapolate that picture to explain the formation of large groups of stars is more

complicated. For example, a complete model of Star Formation should formulate

correctly the necessary rates and efficencies of formation required to populate a galaxy

like ours, but also those of more or less active galaxies. It would also need to be a

sort of general scheme that could explain the formation of stellar populations with

similar characteristics (for example their mass distributions) in completely independent

environments. It would also need to unify the physics relevant to the prime material

(interstellar clouds) and the final product (the stars). Progress has been made, but

nowadays Star Formation, as a global theory, still has many untied knots.

Stars form in molecular clouds, composed mainly of molecular hydrogen, which

are the densest (n> 103 cm3) and coldest (T~ 10 K) components of the Interstellar

Medium. A significant fraction of this molecular material exists in the form of large

complexes called Giant Molecular Clouds (GMCs), with masses of 104_106 M. and

typical sizes of 10-100 pc.) GMCs are usually surrounded by extended envelopes of

atomic Hydrogen with typical masses of 106 Mo.

Practically all known GMCs with distances of less than 3 kpc have been forming

stars during the last 10 million years and we have direct evidence for this assumption:

* First, many HII regions located at the edges of molecular clouds are being expanded by

the winds of young, massive stars. By young and massive we understand O and B spectral










types, with fast nuclear burning rates that result in lifetimes much shorter than the age

of the galaxy. Also, these objects are usually located in groups, called OB Associations

(OBAs). These associations usually have spatial densities below the threshold for Galactic

tidal disruption (Ambartsumian, 1947). This fact provides further evidence -in this case

dynamical- that star formation is recent.

* Second, with the aid of infrared and millimeter-wave detectors developed in the last

three decades, we are able today to see through the optically thick clouds where stars

form -an impossible task for optical telescopes. This way we have been able to observe

stars and even proto-stars while they are still embedded in their parental clouds. These

embedded stellar populations are even younger than OBAs, with typical ages of 1 Myr or

less.

Also, from observations of embedded clusters in nearby GMCs (d < 2 kpc), there

is observational evidence that the majority of the stars in GMCs are formed in clusters

(Lada et al., 1991b; Carpenter, 2000). Moreover, rich clusters (100 members or more)

clearly dominate over small groups, as they contain more than 80% of the observed

embedded stellar population (Porras et al., 2003; Lada & Lada, 2003).

Unfortunately, the dominance of large embedded clusters in available catalogs

might be slightly biased by an incompleteness at the small cluster regimes. Among the

reasons for this are: a) systematic surveys of molecular clouds aiming for the detection

of an embedded population are rather scarce; b) searches for embedded clusters, if any,

are usually limited to those zones with signposts of formation (e.g. the presence of very

luminous infrared sources); c) surveys are mostly based on monotonic wavelength counts,

with poor corrections for background contamination. This way available surveys have

led to the spotting of only the richest clusters. It is only in a few cases when there is a

search -either additional or separate- for low density groups and distributed embedded










populations1 The main reason is that small groups are logically much more difficult to

detect, especially if there are mostly composed of low mass stars (which are fainter), with

large spatial distributions and projected against a high background of reddened sources.

Embedded low mass stars are clearly harder to observe because they are intrinsically

faint. Even so, spectroscopic studies reveal that OB associations have a much larger

number of low mass than massive stars, in proportions that are coincident with the

distribution or Initial Mass Function (IMF) of stars in the field. However, the spatial

density of the low mass component is rarely above that of faint field stars and thus the

exact fraction of low mass stars in young populations is difficult to calculate from stellar

density counts alone (Lada & Kylafis, 1999). Fortunately, if stellar associations are young

enough (3 Myr or less) then low mass stars have circumstellar material that causes them

to have an excess of infrared emission, and this makes them distinguishable from field

stars. These kinds of objects, known as Line-Emission or T Tauri stars, are thus a good

tracer of the low mass component of young stellar populations, but their observation is

subject to uncertainties related to the eventual weakness of line emission, the quality of

the photometry required to observe the excess, and the eventual contamination from field

stars (see Chapter 4).

The process of formation of low mass stars, which leads to their coexistence with

massive stars is also poorly understood. Some existent models of cluster formation are

able to account for the observed spatial distributions of stars in clusters, but fail to match

the observed physical conditions of dense cores where clusters form (Bonnell et al.,

1997), or do not fit the number distribution of the observed IMF (Mouschovias,



1 The term distributed refers to stars for which their formation process cannot be di-
rectly associated to a group or a cluster. For example, it could refer to stars formed in
isolation or stars originally formed in a group but dispersed to the point that the group is
no longer distinguishable (e.g. Li et al., 2002; G6mez et al., 1993; Carpenter, 2000,.)










1991). Also, dense molecular cores are expected to experience significant fragmenta-

tion prior to condensation of proto-stars, a process that is not completely understood

either. The current hypothesis is that marginally stable cores experience cooling via dissi-

pation of magnetohydrodynamic turbulence in highly extinguished cores (Myers, 1998),

which leads to the fragmentation of the core into a matrix of molecular kernels. The

kernels will end up forming stars of different masses via competitive accretion, with the

most massive stars either forming closer to the center of the core where accretion rates

are higher or from initially larger kernels (Bonnell et al., 2001). In this scheme low mass

stars will form preferentially in the outer parts of cores, resulting in a primordial mass

segregation.

The puzzle of the global properties of star formation in GMCs, with a complete

understanding of the mechanisms that lead to the dominion of large groups and the for-

mation and role of low mass stars, can only be solved by studying GMCs in a systematic

approach. This means that entire GMCS should be observed one at a time, with instru-

ments powerful enough to detect low mass stars. Also we need to cover as much area of

the cloud as possible, independently of the presence of rich clusters signposts, so that

low density populations if any, can also be taken into account. In such surveys, we could

obtain unbiased statistics of the embedded stellar populations, and it would be easier to

form a global picture of the stellar birth phenomenon.

1.2 Motivations for the Study of the Rosette Complex

The Rosette Complex is a giant star forming region where a very large OBA,

NGC 2244 has formed. This OBA is evacuating the center of its original cloud by means

of a powerful ionization front created by the winds of its members. At the southeast

edge of this region, there is a large molecular cloud, where several embedded clusters

have been detected. These characteristics make the Rosette an excellent laboratory to

investigate the properties of very young stellar populations. The region has been stud-

ied extensively in terms of its main features, the Rosette Nebula, its OB association










NGC 2244 and the physical properties of the molecular cloud. However, the charac-

teristics of the embedded populations have been studied only to a very superficial level

and it is unknown if there are other clusters, if they share the cloud with a low density

population and more important, what is their relation to the prominent NGC 2244. The

molecular cloud and the nebula appear to be in clear interaction, and a basic question

is how the formation of the new clusters is related to this interaction. One approach to

this problem, for example, would be to study the properties of clusters as a function of

distance to the nebula and see if any significant differences arise, which would be proof of

the influence of the OB association in the new episode of formation occurring in the cloud.

One of the main goals of this thesis is to determine, to the best possible level, the

total number of young stars in the Rosette Complex, as well as their distribution, and

relative properties. The core of the thesis is a new near infrared survey of the region

made with the instrument FLAMINGOS, developed at the University of Florida, which

can detect stars in the Rosette down to the low mass regimes -a task that has not been

accomplished yet. After separating from the catalog the best candidates for young stars, we

apply a technique based on the calculation of local surface densities in order to determine

the location and extent of the known clusters. The selection of stars by their infrared

excess -determined with the use of near-infrared colors, improves the cluster detection

techniques used in single band photometry studies for other clouds.

In the first chapter of this thesis we make a review of the previous studies of the

Rosette Complex region. The review follows a roughly historical line, and ends with the

few embedded population studies done previous to this work, motivating the necessity for

our new observations.

The second chapter of this thesis is dedicated to the description of our Rosette

Complex near-infrared survey, detailing our observations, data reduction methods, and

data quality assessments.










The third chapter describes the analysis applied to the photometric catalogs resultant

from the survey. From this analysis we attempt to improve the discussion about the

distribution of star formation in the RMC.

The fourth chapter describes a complementary millimetric radio wave study of eight

RMC clusters, which has the goal of discussing the interaction between embedded star

clusters and the remnants of their parental cores.

The fifth chapter describes the use of near-infrared colors of stars to create a detailed

extinction map of the Rosette Cloud, which allows us to compare some properties of

the clusters with those of their forming cores. We also include a first approach to the

calculation of the cluster masses, which allows us to study star forming efficiencies in the

complex.

Finally, we present a summary of the results of the thesis and a discussion on future

work.















CHAPTER 2
THE ROSETTE COMPLEX IN MONOCEROS

2.1 Historical Perspective

The Rosette Complex (1=207.0, b=-2.1) is located at the anticenter of the galactic

disk in the constellation of Monoceros. The region is very popular, partly because of the

lIigg.rinfg beauty of its main feature: a very extended emission nebula which hosts a

large central HII region, evacuated by the winds of a central OB association (see Figure

2.1).


Figure 2-1: The Rosette Nebula. Credit: Canada-France-Hawaii Telescope / 2003


The complex is part of a much larger structure known as the Northern Monoceros

Region. This region comprises the Mon OB 1 Cloud (host of NGC 2264 and the Cone








8


Nebula), the Monoceros Loop, and the Mon OB2 Cloud in which the Rosette is one of the

most prominent features (see Figure 2-2).


13 L


11 L


6h 48rh 40m 32m 24m
a (1950)


6h16m


Figure 2-2: The Location of the Rosette Molecular Cloud in the Context of the Mono-
ceros Complex region, from Perez (1991).



The catalog name for the Rosette can be somewhat confusing because it is not

unique: The nebula itself is usually cataloged as NGC 2237 or NGC 2246, (especially

by amateur observers) although NGC 2237 originally referred to the brightest patch at

its west side and NGC 2246 originally pointed to a bright zone at the eastern side. In

addition, while the central cluster is usually known as NGC 2244, it has also been


I I I


NORTHERN MONOCEROS





i i i lql I IT I I










cataloged as NGC 2239. However, this designation historically referred to the brightest

star in the region, 12 Monocerotis.

The cluster was first noticed by Flamsteed in the late 17th century and later reported

by William Herschel -who did not notice the nebulosity- and John Herschel, who

discovered several of its most conspicuous features and reported them in his general

catalog (Herschel, 1864, NGC 2239 = GC 1420).

Other parts of the nebula (NGC 2237 and NGC 2246) were reported by Swift (1886)

who cataloged the object as being "pretty bright [pB], very, very large [vvL] and diffuse

[diff]." Afterwards, the region was formally known as the "Swift Nebula," until the name

"Rosette" became more popular. The total extent of the Rosette was not determined until

the first photographic plates were obtained by Barnard (1894).

Two of the first applications of Rosette Nebula photographic data were made

by Hubble (1922) in his study of diffuse nebulae associated with massive stars, and

Minkowski (1949), who published a photographic plate study along with a first discussion

on the expansion of the HII region by the central cluster O stars and the possible existence

of Bok globules. He estimated the mass of the nebula to be 104 M.and suggested that it

could be "surrounded and probably embedded in obscuring material," thus proposing the

existence of the companion molecular cloud.

2.2 The Rosette Nebula and the Young Cluster NGC 2244

2.2.1 The Rosette Nebula

A tabulation of the different methods used to determine the age of the Rosette

Nebula was done by Ogura & Ishida (1981). These varied from studies of the properties

of the central cavity (Kahn & Menon, 1961; Lasker, 1966) to evolutionary models of the

HII region based on the luminosity of the stars (Hjellming, 1968). Other methods involve

time scales of radiation pressure (Mathews, 1966, 1967), estimates of the formation time

for dark globules in the nebula (Herbig, 1974), and the separation of [OIII] emission lines

(Smith, 1973). The mean value of all these age estimates is approximately 3 1 x 106 yr.










A series of studies by Celnik 1983; 1985; 1986 discussed the global physical

characteristics of the Rosette Complex. The first two of these are dedicated to the nebula,

while the third one is a model of the interaction with the molecular cloud. In the first of

the articles, he presented a map of the Ha emission in the nebula region, and calculated

a total integrated flux density of 5 x 10 11W-m2 within 60' from the center of the HII

cavity. He suggested that the emission is contained in a more or less symmetric ring with

a peak at 16' from the center.

In the second paper, Celnik reported radio continuum observations (1410 and 4750

MHz) from which he was able to determine that the nebula is bound by ionization,

forming a spherical shell with radius of 40 pc (about 85') and a total ionized matter

mass of 2.3 x 104 M Using the H112a and Hel 12a recombination lines (4619 and

4621 Mhz) Celnik calculated a He+ abundance of 0.120.03 and a non-LTE electron

temperature for the nebula of Te = 5800 700 K almost 1100 K above the LTE with

no observable gradient with respect to the radial distance from the center.

However, the investigation by Shipman & Clark (1994) revealed a good fit to

a T oc r-,, a = 0.4 model for the temperature gradient in the nebula cavity which,

interestingly, could not be adjusted to the observed IRAS emission. Instead, they found

that this temperature gradient was better adjusted to a = 0.05 for r < 47' and a = 0.2 for

47' < r < 65'.

2.2.2 NGC 2244

The prominent OB association that is presumed responsible for the evacuation of

the central part of the nebula has been the subject of many interesting studies over the

years. The distance to this young cluster (and therefore to the entire complex) has been

estimated many times with slightly different results. Table 2-1 is a compilation of these

values, from which the most commonly used is 1600 to 1700 pc.

Some of the first visual photometric studies on NGC 2244 were made by Johnson

(1962), who estimated the mean color excess in the cluster to be E(B-V)= 0.46 for









Table 2-1: Distance Estimates to the Rosette (NGC 2244)

Author Value (pc) Method
Johnson (1962) 1660 Photoelectric Photometry
Ogura & Ishida (1981) 1420 Visual Photometry
Perez et.al (1987) 1670 Visual Photometry
Park & Sung (2002) 1660 Visual Photometry
Hensberge et.al (2002) 1390 Spectroscopy


Rv = Av/E(B V) = 3.0. This was confirmed by Turner (1976) and later Ogura & Ishida

(1981), who suggested a value of R = 3.2 0.15. Ogura & Ishida (1981) also pro-

posed an age of 41 Myr and a star formation efficiency of 22% for the cluster. Later,

Marschall et al. (1982) completed a proper motion study of 287 stars in the NGC

2244 area. They confirmed membership for 113 objects, 52 of them from the list of

Ogura & Ishida (1981).

A study that combined photometry as well as spectroscopy was completed by

Perez et al. (1987). They found that some members of NGC 2244 presented anomalous

values of R, possibly suggesting the coexistence of main sequence stars with very young

objects -likely T Tauri stars. This was confirmed with a uvbyp photometry study by

Perez et al. (1989), in which 4 members presented evidence of being true pre-main

sequence (PMS) objects. They also confirmed the age of NGC 2244 to be below 4 Myr

but spread towards younger values, thus confirming a model of continuous formation.

A study of great importance was performed by Park & Sung (2002). They obtained

UBVI and Ho photometry for the cluster. They were able to determine membership for

a total of 30 cluster sources and to extend the list of known PMS candidates to 21. They

subsequently identified members coincident with ROSAT point sources catalogs and

spectral types from Verschueren (1991) (see section 2.2.3). Six of the PMS candidates

were confirmed as X-ray sources. In Figure 2-3, we show the Park & Sung (2002)

relation between alpha emission and V-I color for NGC 2244. (the relationship actually










would hold for any optical or infrared color). In this figure, PMS stars are clearly located

above the main sequence.

Later, using evolutionary models Park & Sung (2002) showed that most of the PMS

stars and PMS candidates in their sample appear to have masses close to 1M. and an

approximate mean age of 0.4 to 0.9 Myr. Because they also estimated the main sequence

turn-off age of the cluster to be 1.9 Myr, they showed that the cluster has not stopped

forming stars yet.

Another important calculation in this article is the Initial Mass Function (IMF) of

NGC 2244. They found it has a flat (F=-0.7) IMF slope in the range 0.5 < logm < 2.0. By

comparing directly to the IMF model of Scalo and to the observed mass function of

NGC 2264, they demonstrated that NGC 2244 is highly dominated by massive stars, thus

confirming its status as a giant OB association.

2.2.3 Spectroscopic Studies

The most complete spectroscopic study of NGC 2244 was done by Verschueren

(1991), (see section 2.2.2) and it has been widely used in the literature. In particular,

Park & Sung (2002) used data from this study to identify the spectral types of candidate T

Tauri stars in NGC 2244.

A low resolution, single slit investigation by Hensberge et al. (1998) of 2 members

and 3 field stars in the region of NGC 2244 yielded evidence that these were chemically

peculiar, possibly magnetic stars. Later, Hensberge et al. (2000) performed spectroscopic

analysis of the binary member V578 Mon, which resulted in an estimated distance

slightly lower than other photometric estimates (see Table 2-1). They also calculated the

age of the system to be 2.30.2 Myr.

Finally, Li et al. (2002) presented low resolution spectra for a sample of X-ray

counterparts from the ROSAT PSPC survey (see also Gregorio-Hetem et al. (1998)). They

were able to confirm that five sources had strong Hc emission. Two of the stars were












-2 A


Ha-


A

0A -- &.._





0 1 2 3
V-I

Figure 2-3: A Ha vs. V-I diagram for NGC 2244 from Park & Sung (2002). The
solid line represents a ZAMS relation while the dashed line is a selection
limit. Filled triangles are PMS stars while open triangles are PMS candi-
dates. Bright members are marked with dark filled circles. X symbols are
X-ray sources and dots are non-members.

found to be Herbig Ae/Be and two others had WTTS profiles. These data indicate that

X-rays are an efficient tracer of young populations.

2.2.4 Near Infrared Studies

Recent surveys in the near-infrared permit investigation of the extension and

structure of the cluster. In the study by Li (2005) they analyzed data from the 2MASS all

sky survey, and suggested that NGC 2244 had a second component located approximately

6.6 pc west of the core center. Data from the FLAMINGOS survey reported in this thesis,

appear to confirm the existence of this second association (see 4), which is coincident

with the area originally labeled as NGC 2237. This area is particularly interesting because

it contains large dust structures known as "elephant trunks", as well as other types of very

young condensations of material which suggest very recent formation.










2.2.5 X-ray Studies

NGC 2244 is an important target for X-ray studies due to the interest in investigating

the nature of massive stars as sources of high energy photons. The ROSAT Consortium

observations yielded 34 X-ray sources in NGC 2244, with typical energies of 1030 1032

ergs s 1. Six of these X-ray sources are PMS candidates as reported by Park & Sung

(2002). Also, Berghifer & Christian (2002) studied NGC 2244 ROSAT sources and

found that objects with the faintest X-ray emission have very high X-ray to optical

luminosity ratios. They noted that the number of X-ray emitters associated with Ha

emission in NGC 2244 is remarkable. Taken together, these data give strength to the

hypothesis that many X-ray emitters are young late type stars.

2.3 The Rosette Molecular Cloud: Structure

2.3.1 CO studies

A substantial part of the Interstellar Medium (ISM) exists in molecular form. Molec-

ular hydrogen (H2) is stable and abundant, but unfortunately is not easily detectable

because H2has no permanent dipole moment and therefore its transition probabilities are

very small. CO is considered instead the best tracer of molecular gas because of its high

and constant abundance in molecular hydrogen clouds.

First attempts to detect CO emission associated with the Rosette nebula were

unsuccessful as they pointed at the nebula region, which is mostly composed of neutral

and ionized gas.

The observations reported by Blitz & Thaddeus (1980), which targeted the southeast

adjacent region of the nebula, were the first successful detections of molecular gas in the

Rosette. Their NRAO survey mapped over 80% of the 12CO emission in the area of the

cloud with a 1' beam size, and yielded information about its large scale distribution. They

estimated the angular extent of the cloud to be 3.5deg (98 pc at a distance of 1600 pc)

and labeled the most prominent sub-structures. In the adapted map of Figure 2-4 we

show an optical picture of the Rosette nebula from the DSS overlayed with contours of









12CO from the Bell Labs maps of Blitz & Stark (1986). We included the nomenclature of

Blitz & Thaddeus (1980).

Particularly important regions are the Monoceros Ridge (region A1-2) which is

literally a region of gas compression at the cloud-nebula interface; the Cloud Central

Core (Al-1), which hosts the most massive clumps in the cloud and is the strongest

region of star formation; the cores D and G, which are separated from the main body of

the cloud but have ongoing star formation; the IRS core, which hosts the massive proto-

binary AFGL-961 (see 4 and 5.1); the Back Core B, which is more loose in structure

than the regions near the nebula; and the Arm or E core, which despite its brightness

contains no significant star formation (no IRAS sources, or near-infrared clusters have

been found in this core so far).

In a subsequent study Blitz & Stark (1986) mapped the 12CO and 13CO emission

with improved sensitivity at the AT&T Bell Labs, uncovering the high degree of clumpi-

ness of the cloud. The study of Williams et al. (1994) also made use of the data from

Blitz & Stark (1986) and listed a total of 95 clumps. The clumps with evidence of star

formation had larger peak temperatures, larger densities and also were more gravitation-

ally bound compared to clumps from the Maddalena Complex, a cloud with very low star

formation. Later, Williams et al. (1995) showed that about half of the clumps in the RMC

were gravitationally bound and the rest were supported by pressure from the interclump

medium, which was shown to be mostly atomic and about 40 times less dense. In figure

2-5 we show the locations and relative sizes of the clumps from Williams et al. (1995).

From the clump central velocities Williams et al. (1995) found that the cloud has a

well defined velocity gradient of about 0.08 km.s- 1pc1. Also, the negative correlation

between clump mass and clump to clump velocity dispersion suggested that the system is

still far from equipartition even though it is dynamically evolved. Williams et al. (1995)

also found the star forming activity to be more intense in the ridge and the central core

areas, near the interface cloud-nebula: clumps located near the nebula presented larger











ROSETTE COMPLEX
6.07297
IRAS 25 micron. Digital Sky Survey



5.47608 -









0 P
4.87919 -




U .
04.28231 -





3.68542 -


AT&T Bell 12CO Levels

3.08853 I I II
6.65547 6.61515 6.57483 6.53451 6.49419 6.45387
Right Acension (J2000)


Figure 2-4: A map of the Rosette Complex area. The background image is a DSS plate of
the IRS survey at 25 microns. The contours represent 12CO integrated inten-
sity levels from the survey of Blitz & Stark (1986). Indicated with labels are
the main regions of the molecular complex identified by Blitz & Thaddeus
(1980)


excitation temperatures, average densities and star forming efficiencies and could be

translated as brought clues of evolution. Other properties of the clumps (mass, sizes or line

widths) did not show any significant variations along the cloud.

Another CO study was done by Schneider et al. (1998a). The observations focused

on the central part of the cloud, detailing the structure of the midplane star forming

cores. They paid special attention to the IRS core, where the source AFGL-961 is located














2I
'' d:_ 1; ,. <* -_


-2






208.5 208 207.5 207 206.5 206 205.5
Galactic Longitude ()

Figure 2-5: Locations and relative sizes of clumps in the Rosette Molecular Cloud from
Williams et al. (1995). The size of the symbol is proportional to the mass of
the clump.


(see also section 5.1) and pointed out the ample blue wing emission due to the powerful

outflow from this object. In a complementary study (Schneider et al., 1998b), examined

the CII emission (158pm) at the ridge, the central core and the IRS core. They found

weak but significant C+ emission deep into the molecular cloud cores and suggested that

the distribution agrees well with a clumpy molecular cloud exposed to a low level UV

radiation field. The penetration of UV photons in the cloud is apparently facilitated by a

high density contrast clump-interclump medium.

The clump mass spectrum in the RMC has the form dN/dM oc M where x =1.6,

with small variations in the exponent depending on the range, bin size and beam reso-

lution used (for example Williams et al. suggest that x is closer to 1.3). The exponent

in this power law is similar to other clouds (Blitz, 1993), but what is more important, it

shows that the clump mass spectrum is much shallower than the observed stellar IMF

(x = 2.35). This shows that although small clumps have a larger number proportion, most

of the mass is contained in only a few big clumps, while for stars both numbers and total

mass are dominated by the lowest mass bins. Interestingly enough, the power law index










is in fact similar to that corresponding to the mass distribution function for embedded

clusters (Lada & Lada, 2003), which is suggestive of a uniform star formation efficiency

for most star forming cores.

The more recent survey of Heyer et al. (2005), obtained with the wide field array

SEQUOIA at the FCRAO 14m telescope have a resolution of 45" at 115 GHz and 47" at

110 GHz. The maps reveal "textural variations" in the 12CO emission across the complex,

with a brighter emission component within the nebula projected radius (approx. 40

pc from the center as defined by Celnik (1986)) and weaker, more extended emission

outside this ionization edge. HBW05 suggest that the weaker emission is probably due

to subthermally excited material with lower densities. They also calculated the total

molecular mass of the cloud to be 1.6 x 105M. from 12CO, and found a LTE mass of

1.16 x 105M. from 13CO. Moreover, they were able to apply a Principal Component

Analysis (Heyer & Schloerb, 1997) to determine the turbulent flows and the turbulence

scale in the RMC. This analysis reveals more significant variations in the velocity

structure of the cloud at the regions located within the ionization than in the more

diffuse, external component. This fact reveals the interaction of the cloud and the HII

region. They suggested, however, that these interactions are still very localized, and have

not affected the global dynamics of the cloud yet.

2.3.2 Interaction with the Rosette nebula

In his third study of the Rosette, Celnik (1986) focused on comparing his Ha map

and radio continuum observations of the nebula (see section 2.1) with the CO map of the

molecular cloud from Blitz & Thaddeus (1980). Celnik constructed a complex model of

the distribution of the main CO cores (see Fig 2-4) in the context of the HII region and

estimated the rotation center of the cloud at (u, 8) = (98.1615,4.3287,J2000.0). Finally,

he re-calculated of the mass of the entire complex by adding the total mass of ionized

atoms, stars, dust and molecular gas, resulting in 3.3 x 105 M..










Cox et al. (1990), used the available IRAS data (12, 25, 60 and 100pm), and

determined in great detail the distribution of dust and compared this to the distributions

of ionized and molecular gas. Additionally, they were able to estimate a total infrared

luminosity of roughly 1.1 x 106 L. or about 50% of the available luminosity from

the cluster NGC 2244. Warm dust (usually present near an OB association) typically

emits strongly at the four IRAS bands. However, Cox et al. also showed that in the

Rosette, while the 60 and 100pm emission were quite strong at regions of ionized and

neutral gas (nebula), the 12pm emission was preferently located beyond the limits of

the ionization front (molecular cloud), suggesting a heavy rate of destruction of dust

grains from UV radiation from the cluster. Surprisingly, the 25pm emission was found

to be significant in some parts of the ionized nebula, possibly due to the existence

of a second type of dust particle that is more resistant to UV photons. This was also

suggested by Shipman & Clark (1994), who found that the maximum temperature in

the shallow temperature gradient found in the nebula (see section 2.1) seems too low

to sublimate ice mantles in grains and too low for grains to emit significantly in 12

or 25pm -a problem possibly solved with a second type of grain in the region. Later,

Shipman & Carey (1996) suggested that line emission could be contributing strongly to

the IR emission of the nebula, and suggested, once more that the presence of a "hot dust"

component is necessary to model this emission, especially for the 25pm.

Figure 2-6 shows the superposition of the 12pm emission from the IRAS survey

and a 1400 MHz radio continuum emission map from Holdaway, Braun & Liszt (un-

published). The infrared contours indicate that the warm dust emission defines a shell

that encloses the ionization front, showing the effect of heavy dust destruction by the

nebula. The overposition of these maps defines very clearly the region where the HII

region impacts the molecular cloud.

Kuchar & Bania (1993) made a complete map of HI emission at 21cm using the

Arrecibo telescope. They found that atomic gas in the Rosette Complex is distributed














ROSETTE COMPLEX


image: L-band 1400 Mhz
- IRAS 12/.


5.5



o
O
0

oL
5.0






4.5






4.0


98.0 97.5 97.0 96.5
RA (J1950.0)


Figure 2-6: IRAS 12pm emission map superimposed on a 1400 MHz radio continuum
emission map by Holdaway, Braun and Liszt, NRAO.


in three main regions which form a rough, extended shell around the optical nebula

and beyond the molecular cloud, with a center of expansion at (a, 8) =(97.95,4.97,

J2000). This shell (according to their calculations) would have a mass close to 2x 104

M which implies a budget of kinetic energy for the shell expansion of approx. 4x 1048

ergs, or 2% of the total energy available from the stars in NGC 2244.


'I I
r
F'I
i;










High Energy studies

The possibility of interaction between the HII region and the star forming cloud, as

well as the location of the Rosette Nebula near the edge of the Monoceros Loop -also

known as the Monoceros Supernova Remnant- (Davies, 1963), has motivated a number

of studies aimed at investigating the high energy photon emission in the interaction

regions.

Deep HO+[NII] photographic plates by Davies et al. (1978) suggested a correlation

between a filamentary structure observable in Ha emission and a Rosette nebula feature

observable in 240 MHz radio waves. This feature was proposed as evidence of loop-

nebular interaction and confirmed by decameter (Odegard, 1986) and diffuse X-ray

emission (Leahy et al., 1986) observations. Later, high energy (100 MeV) y-ray images

from EGRET (Jaffe et al., 1997), revealed a feature partly coincident with the filaments

and apparently significant (7o) over expected diffuse emission. If real, these y energy

photons would be a product of the interaction of charged particles with the dense ambient

medium at the shock region. Recently, the HEGRA system of atmospheric Cerenkov

telescopes at IAC was used to calculate the cosmic ray emission from the loop-nebula

interaction region, but no significant TeV energies were found (Aharonian et al., 2004).

When the Rosette Molecular Cloud was confirmed as a region of star formation,

Gregorio-Hetem et al. (1998) used ROSAT data again, this time to map the MonR2

cluster and the Rosette Molecular Cloud areas in order to confirm a correlation between

star forming cores and clusters of X-rays sources. They found strong X-ray emission in

NGC 2244, the ridge of the cloud (A1-2 in fig 2-4), and at the cloud core area (Al-1),

but the resolution was poor and individual sources could not be resolved. They suggested

that molecular cores known to have active star formation but failing to show significant

X-ray emission, could be predominantly forming low-mass stars. They also suggested that

detectable X-ray counterparts are in most cases Herbig/AeBe or T Tauri stars, as found in

NGC 2244 (Li et al., 2002).










The more recent observations of the Rosette Complex done with Chandra

(Townsley et al., 2003) have resolutions of only a few arcseconds, thus allowing for

the detection of X-ray counterparts for 75% of the OB members of NGC 2244. One of

the most interesting results of this study was the confirmation of a second, soft diffuse

emission which probably originates from the O star winds and is later brought to thermal-

ization by wind-wind interactions or by the shock with the surroundings, in this case the

molecular cloud (see figure 2-7). This X-ray plasma surrounds the OB association and

fills the nebula cavity completely.


6:3f00 W. .. 34:p0 3p 33:p0 3 32:p0 31:p0 .
Right Ascension (J2000)

Figure 2-7: A 0.5-2 keV Chandra image of the Rosette Complex. The emission has been
smoothed to highlight the soft diffuse emission that originates in the neb-
ula and propagates into the molecular cloud. Credit: Townsley et al. and
NASA/Chandra X-Ray Observatory (2003).










2.4 The Rosette Molecular Cloud: Embedded Populations

The coincidence of massive clumps and luminous IRAS sources pointed out by

the study of Williams et al. (1995) strongly suggested that star formation had already

taken place across the molecular cloud. However, the poor spatial resolution of the

IRAS point source survey did not allow the resolution of individual members of an

embedded population. Early near-infrared studies (e.g. Perez et al., 1987) did not cover

the molecular cloud areas, and optical photometric studies were incapable of detecting

obscured populations.

An exploratory near-infrared survey (JHK) by Phelps & Lada (1997) that made use

of the imager SQIID finally confirmed the existence of embedded clusters in some of the

most massive clumps from the list of Williams et al. (1995) that were associated with an

IRAS source. They were able to distinguish seven deeply embedded clusters with bright

nebulosities, and suggested that clumps not forming a cluster, might not be physically

bound. The location of the seven Phelps & Lada (1997) clusters is shown in Figure 2-8

Complete area coverage of the Rosette Complex in the near-infrared was first

accomplished with the release of the All-Sky 2MASS survey catalogs. The 2MASS

survey was a major gain in data uniformity but unfortunately not in sensitivity. Due to

the distance to the Rosette (1.6 kpc), the 2MASS completeness limit (K=14.3 mag) is not

deep enough to study the low mass end of the IMF.

So, what is the next logical step in the study of the Rosette Complex? The existence

of embedded clusters in the Molecular Cloud means that the cloud is actively forming

stars and that at least a fraction of the new stars were formed in clusters from the collapse

of some of the most massive clumps of molecular gas. This leads to two problems of

importance:

2.4.1 Dominance of Cluster Formation in the Rosette Complex

The first problem is to determine if star formation in the RMC leads to a dominance

of rich clusters. Are there any low density groups as well? Is there any evidence for a
























Galctic Longitude (1)
Figure 2-8: The location of the clusters identified in the study of Phelps & Lada





2.0

-2.2 .
-5I

onurvey Bell Lobsntours indicate CO emission from the maps of Blitz & Stark
-2.4 +Embedded Clusters P

208.0 207.5 207.0 206.5 206.0
Golactic Longitude (M)


Figure 2-8: The location of the clusters identified in the study of Phelps & Lada
(1997). The background image is an optical plate from the Digital Sky
Survey. The contours indicate 12C0 emission from the maps of Blitz & Stark
(1986)



distributed population? For example,Carpenter (2000) showed that the MonR2 region,

might be harboring a low density population, counting for up to 9% of the total number

of young stars. The nature of such low spatial density members is not clear, as it could

be formed independently of the cluster population, but also could be the result of the

dispersal of an older high spatial density population. Another study that attempts to

account for the contribution of distributed populations in star forming clouds was done

by Li et al. (1997), who found that for the L1630 cloud, where Lada et al. (1991b)

found unequivocal dominance of cluster formation. The fraction of infrared excess stars

in the inter-cluster areas of the cloud was found to be very small, suggesting that the

contribution of low-density formation was almost negligible.










We need to consider that the RMC is located 4 times further away than the well

studied Orion or Perseus molecular clouds (d = 300 500 pc), where reasonably deep

observations can easily detect low mass stars. Equivalent detections in the RMC would

need observations at least 3 magnitudes deeper. Furthermore, the RMC is located in the

direction of the galactic anticenter (1=207deg) and closer to the galactic disk (b=-2)

than Orion (b=-16.3) or Perseus (b=-20.6). As a result, the density of the field population

towards the RC is very high, and any faint, low mass stars are probably well mixed

with foreground and background sources. These problems would make very difficult to

detect other clusters or a low density population by means of single band stellar density

counts, as in other cloud surveys (Lada et al., 1991b; Carpenter, 2000).Also, depth limited

databases like 2MASS are not sensitive enough to study the Rosette Complex. Deep

multi-wavelength photometry, capable of rendering infrared colors, for even low mass

populations is necessary to separate members from field stars by means of accurate

extinction statistics and counts of infrared excess stars.

2.4.2 The Hypothesis of Sequential Star Formation

The second problem to be understood, relates to the physical processes that led

to the formation of stellar clusters in the Rosette: Are those processes similar to those

occurring in other star forming clouds?. The current hypothesis is that the formation of

star clusters in the RMC was possibly stimulated by the interaction of the HII region

and the cloud. The expansion of the Nebula via the ionization front generated by the

strong stellar winds of the massive association NGC 2244, results in a shock front which

interacts with the gas of the molecular cloud, as shown in some of the studies mentioned

above. The hypothesis is that the shock front directly stimulated the collapse of clumps

which then formed the clusters. This model is known as sequential star formation (SSF),

and was developed theoretically by (Elmegreen & Lada, 1977).

In the study of Williams et al. (1995), it was shown that the cloud had larger values

of excitation temperature, clump density and possibly, star formation efficiency near the










HII region. However, there are not significant differences among characteristics of cluster

forming clumps, namely mass, size or line width across the cloud. Could this mean

that other massive clumps, either those not included in the areas of the Phelps & Lada

(1997) survey, or those not associated with a luminous IRAS point source could also

have formed stars recently, even if their location is not favorable with respect to the

shock front? In other words, how feasible is the hypothesis of SSF? The detection of

additional embedded populations would allow us to determine for once if star formation

is preferentially located near the shock front of the nebula expansion. We might also

be able to find a relation between the characteristics of the embedded clusters and their

distance to the HII region that could support or discard the hypothesis of SSF.















CHAPTER 3
A NEAR-IR SURVEY OF THE ROSETTE COMPLEX: OBSERVATIONS

3.1 The FLAMINGOS GMC Survey

As we discussed in Chapter 1, a systematic and thorough investigation of the young

star population of GMCs is the key to understanding the global aspects of the problem

of Star Formation. Historically, the large distances and large angular sizes of GMCs,

made it difficult and costly to perform surveys of embedded populations which could

render both photometric depth and area coverage. Technological limitations were also a

factor, with infrared detectors having very small areas: until the early 1990s, near-infrared

arrays were no larger than 256 x 256 pixels which resulted in rather poor resolution and

a small field of view (FOV). For example, the survey of the region L1630 in the Orion

Molecular Cloud by Lada et al. (1991b) used a 58 x 62 pixel device which rendered a

FOV of only 1' x 1', and thus required of 2800 images to cover an area of approx. 0.7

square degrees in the K band. Large devices were developed then, with the instrument

SQIID (Ellis et al., 1993) being the first versatile instrument to allow a high resolution

and a large FOV (1024 x 1024 InSb device with simultaneous quadrant detection in J,H

and K). This camera was used for the first time to survey large areas of star formation

regions in multi-band mode, like rho Ophiuchi (Barsony et al., 1997) and the Rosette

Molecular Cloud (PL97).

Near the end of the decade, the first 2048 x 2048 HgCdTe devices for the use

in astronomical instrumentation were developed (Kozlowski et al., 1998), opening

even better possibilities. The instrument FLAMINGOS (Elston, 1998), developed at

the University of Flhtida., takes advantage of the 4 million pixel detectors by being

designed as a combination wide field near-IR imager and multi-object spectrometer. The

camera has a Lyot stop wheel with a number of stops customized to receive different










input beams slower than f/7 and therefore can provide a wide range of pixel scales for

imaging. For example, on the Kitt Peak 2.1m telescope it renders 0.606" pixels and a

20'x 20' arcminute FOV. This particular setup makes FLAMINGOS an excellent survey

imager as entire square degree areas can be surveyed with a few observed fields. The

instrument has a suite of four filters: J, H, K and K, which cover the whole near-IR

wavelength range from 1.6 to 2.2pm.

Intended as one of the first large scale applications of the instrument, the NOAO

survey program Toward a Complete Near-Infrared Spectroscopic and Imaging Survey

of Giant Molecular Clouds (PI E. A. Lada) is dedicated to the global study of several

giant molecular clouds using FLAMINGOS. One of the two main goals of the survey

is to do a complete imaging coverage in J, H and K of comprehensive areas of the

clouds with a photometric depth that assures coverage down to the Hydrogen Burning

Limit (HBL). Four important GMCs were selected for this survey: Orion B, Perseus,

Monoceros, Cepheus, Serpens and the Rosette.

Observations for the survey program have been carried out over the course of 6

winter observing seasons from 2000 to 2005. The survey was done at the 2.1 and 4.0m

telescopes of the Kitt Peak National Observatory, where FLAMINGOS is a commissioned

instrument.

Although FLAMINGOS is suitable for multi-object spectroscopy (MOS) and

imaging mode in both telescopes, we preferentially performed imaging at the 2.1m

telescope, where FLAMINGOS has a larger FOV, while the 4.0m telescope has been used

essentially for MOS. The imaging observations for the GMCs targets were carried out

iteratively, with Orion B and the Rosette being the first clouds to be completed. After

reduction of the first batch of observations, a first quality assessment was performed by

members of the team and collaborators at the Center for Astrophysics in Cambridge,

Massachussetts during the fall of 2003. Those fields that yielded poor results were

assigned for re-observation during the winters of 2003 and 2004.










In the particular case of the Rosette Complex, a total of 22 FLAMINGOS fields were

observed during the winters of 2001 to 2004 in the 3 available filters J, H and K, resulting

in a very complete coverage of the Rosette Nebula the Rosette Molecular Cloud areas. We

selected the area of the Complex to be covered from the 12CO and 13CO emission maps

of Blitz & Stark (1986) and the 25pm emission map from the IRAS survey. Twenty of our

22 fields are adjacent, while 3 of them (areas 4, G1 and G2 ) were added to enhance the

quality of the observations in some particularly interesting regions. In order to account

for the field contamination, two control fields were observed at close distance from the

survey areas but away from the main molecular cloud emission. The control fields were

observed with an equivalent method to the main survey fields, and have the same depth

as any of our on-source fields. A map showing the positions of the observed fields in the

context of the molecular gas emission (12CO) and the 25 micron IRAS flux in the area

can be seen in Fig 3-1.

For each field we aimed for a total of 1000 sec. on source integration in each

filter (for some fields, weather conditions and defective frames kept us a tad below

this goal), which was done by obtaining a number of short, dithered exposures. For

the J and H filters, we used dither exposures of 60 sec. each each dither, and for the K

band, with a higher sensitivity, we used 20 or 30 sec. dithers, depending on the weather

conditions. Details of the observations, including dates, total integration times, average

seeing, and airmass for each field can be consulted in Appendix A.

3.2 Data Reduction

3.2.1 The Data Reduction Pipeline: LongLegs

Each FLAMINGOS individual image is stored as a FITS file with a size of 16

Megabytes. Each field is observed in three filters and requires combining groups of

dithered pointing. The resultant amount of data for the survey is therefore very large,

and required the development of automated processing pipelines for reduction and

photometry.













ROSETTE COMPLEX
5.67218
IRAS 251 Digital Sky Survey
T--
14 15 :i



4.96886- 10 2

21 3

.
4.26554 1

6


SI20
3.56222- 2





CF2
2.85891 -


AT&T Bell '2CO Levels
+ Rosette Embedded Clusters

2.15561 I I I I
99.5010 98.8144 98.1278 97.4412 96.7547 96.0681
Right Acension (J2000)
FLAMINGOS Giant Molecular Cloud Survey


Figure 3-1: Scheme of the University of Florida/NOAO Rosette Molecular Cloud
Survey. The boxes delimit individual FLAMINGOS fields (20 x 20' after
trimming) over an image of the IRAS 25pm emission in the region. The
labels at the left side of each box refer hereafter to the fields detailed in Ap-
pendix A and the text. Light solid contours represent the extension of the
Rosette Molecular Cloud in CO emission from the survey of Blitz & Stark
(1986). Crosses mark the centers of known embedded clusters from the
previous study of Phelps & Lada (1997).


Our image reduction pipeline, nicknamed LongLegs and programmed by the author,


is a standard IRAF routine script divided into three main phases:










During the first phase, the pipeline rejects defective images and removes bad

pixels. Then, it applies a 3rd degree polynomial linearization correction for every image

on a pixel by pixel basis (IRAF routine IRLINCOR). Dark and flat field data groups are

combined into master flat and dark fields, which are then used to create bad pixel masks.

The second phase of LongLegs is a two-step preparation of pre-combined

data. The algorithm combines groups of 8 adjacent images to create a local sky. Then,

after each data frame has been sky-subtracted and divided by the normalized flat field,

the program reconstructs the individual images dither pattern based on the positions of

the 200 brightest sources in each frame. The program does a first combination of data,

extracts the positions of sources with fluxes larger than a pre-selected sigma level, and

masks them out from individual images to create a new set of "starless" local sky frames.

These are used for a second pass of sky subtraction, flat field division and shift-and-add

combination. The final result of this phase is a set of precombined frames and a first

combined image with analysis quality that only lacks a geometric distortion correction.

On the third phase the pipeline program corrects for geometric distortion, using

a sixth order Chebyshev polynomial solution map constructed from the positional

distortions of a 20x20 pinhole grid mask that is pre-imaged each time the instrument is

corrected internally (corrections indicated slight variations in the geometric distortion

from season to season). The pre-combined data is also re-sampled to half-size pixels,

the dithers are centroid corrected, and re-combined into a final image that is 4096 x 4096

pixels in size and is ready for the photometry pipeline.

3.2.2 The Photometry and Astrometry Pipeline: PinkPack

Our photometry and astrometry pipeline, nicknamed PinkPack and programmed

by Joanna Levine (Levine, 2006), performs stellar profile fitting (also known as Point

Spread Function or PSF fitting) photometry on a LongLegs final product. The script

also uses standard IRAF-DAOPHOT tasks (Stetson, 1987), except for the detection,

which is performed using the S-extractor algorithm (Bertin & Arnouts, 1996). A full










description of Pinkpack can be found in Levine's PhD thesis. We will only mention that

the pipeline gives out a full photometric calibration and an astrometric solution with

respect to the 2MASS All Source Catalog Release data. Calibration of data is done in

the range K = 11 to 14.5 mag. Once a photometric catalog is obtained and an astrometric

solution is calculated, this pipeline combines the data from different filters into a final

merged catalog that contains, for each object, an ID, final RA-DEC coordinates, pixel

positions (in the K band image), and photometry for all the bands, including profile fitting

uncertainties.

The photometry pipeline also has the option of creating and subtracting a median

value image for a specific field in order to enhance the detection of sources in regions

with bright nebulosities. We used this option in all frames that contained bright nebulosi-

ties, although doing the same in non-nebulous regions had no effect whatsoever on the

number of detections obtained.

3.3 Completeness of Sample

In order to estimate the completeness of our sample, we performed intensive

artificial star experiments in 3 selected fields of the survey, each one considered to be

characteristic of a type of region: crowded (with low extinction), sparse (high extinction,

no nebulosity) and with bright nebulosity. Our main goal was to determine mean values of

completeness to apply to the entire survey for our statistical purposes.

The completeness limits for the region containing bright nebulosity emission were

not affected in a greater way than in zones of higher stellar density, although in both cases

the experiments performed slightly better in the sparse fields (see Figures 3-2 and 3-3).

For all the fields, the artificial stars were added partially, in consecutive annuli of 250

pixels from the center of the frame. For each annuli, 100 artificial images with 50 stars in

an uniform distribution were created based on the resultant PSF profiles from Pinkpack

for that specific field. Their magnitudes were adjusted according to the mean zero point

value calculated from the 2MASS calibration. The resultant images were then reduced











Rosette Mo ecular Cloud Artificial Star Experiments (FILTER)

9-20





S16

0 500 1000 1500 2000


20

19 I Crowded
1 8 I0----- -. ,__ Nebulous
-cc
S17 Sparse 1

16 Sparse 2
15 I
0 500 1000 1500 2000


20
19 1 9
L 1 8 -

S17

16

0 500 1000 1500 2000




Figure 3-2: Results of the artificial star experiments described in section 2.4. a) the left
side panels show the turnoff magnitudes (limits of 90% object recovery) by
filter. The average values in this graphs were used as our general complete-
ness limits for the survey.


with the same set of parameters as the original frame, and the positions and magnitudes

of the artificial stars were recovered using the XYXYMATCH routine from IRAF. The

stars in the recovery catalogs were divided by brightness in bins of 0.25 mag, and the

completeness limit was calculated as the bin at which the recovery fraction descended

below 90 percent.











Rosette Molecular Cloud Artificial Star Experiments (AREA)


s-"-
S20 I
19 .. .................... .... I I
.: ....... ... .



16

S0 500 1000 1500 2000



......................... ..... .. J b n
9 1 ...... J bond
18 ------ -cc-H bond
The resultant average completeness limits are K=17.25, H=18. and J=18.50 mag
S 17 I I I
SI I I .... K band


S0 500 1000 1500 2000

























within the limits of acceptable focus quality of the images (see section 3.4); these results
< 19 .. . .I ...
a_ -- ----- --- *---

E 17 I.
o 1 6
E I I I IL
0 500 1000 1500 2000




Figure 3-3: Results of the artificial star experiments separated by type of field. We de-
tected a subtle variation of the recovery limits in the case of too crowded or
too nebulous fields.


The resultant average completeness limits are K=17.25, H=18.00 and J=18.50 mag

within the limits of acceptable focus quality of the images (see section 3.4); these results

rapidly degrade in the areas of high optical distortion. However, these stand for now as

some of the deepest observations of the region, going about 3 magnitudes fainter than

2MASS, and thus assuring the detection of stars around and below the HBL.










3.4 Positional Correction of Photometry

During the assessment of data quality it was noticed that our pipelines had difficulties

adjusting correctly the PSF in certain areas of the chip. The problem was worse towards

the corners of the images, where the stellar profiles were in some cases clearly aberrated

and even presented prominent comas. It is known that the parabolic shame of the primary

mirror has an effect on large detectors, which can be usually corrected with a second

degree surface variation of the PSF profile, but apparently the distortions we observed

had a different origin, because the distortions are not symmetrical, i.e., the four comers

of the images are affected differently. The distortions can also worsen with poor focusing

and bad weather (i.e. mediocre seeing values). One hypothesis based on optical path

simulations (Eikenberry, S. Univ. of Florida, personal communication) is that the

alignment between the primary and secondary mirrors of the KPNO-2.lm telescope

has lost accuracy along the years, affecting the symmetry of the focus and shifting the

center of optical alignment. This defect is unfortunately enhanced by the large FOV of

FLAMINGOS.

The size of the affected area varied slightly from season to season. The area with

minimal distortion is nearly circular, with a center that falls systematically on the pixel

position (3000,2400) for observations made before the fall 2004, and on the position

(3200,2170) for more recent observations. The radius of this area within which the PSF

values stay uniform is variable, with an average of 3200600 pixels depending on the

observing conditions, mainly seeing value, which is affected respectively by the airmass

and the weather conditions at the observatory.

For most of our fields, about 75-95% of the area of the detector contained minimal

distortion, with reasonably smooth PSF profiles and small (< 0.025 mag) photometric

differences with respect to the 2MASS catalogs. Outside this area, the optical distortion

increases quickly, and therefore the shape of the stars and the PSF profiles degraded

to the point that stars presented noticeable aberration comas and larger PSF FWHM










values,especially at the two eastward comers of the detector, which resulted in poor

fittings to the average PSF profile generated from good quality stars in the acceptable

area, and generated a net flux loss with respect to 2MASS that raises sharply to 0.5 mag

in the bad psf fitting areas, independently of the filter.

In terms of the net output to our photometric catalogs, this effect resulted in a

variation of the photometric calibration zero point (ZPT) value across the images, and this

affected the uniformity of the survey from field to field.

In order to correct for this effect, we applied a 6th order Legendre polynomial

correction of the ZPT values as a function of the radial position with respect to the pixel

centers of the optical distortion circles. The solution was developed and constructed

as interactive software by Andrea Stolte. The correction is applied done by fitting a

polynomial to a fiducial line made by the median values of the 2MASS vs FLAMINGOS

differences in radial bins of 300 pixels from the optical distortion center. The correction

was calculated within the ranges 10.0 to 14.0, 10.0 to 15.0 and 10.0 to 16.0 mag in K, H

and J only due to the limitations of sensitivity of the 2MASS catalogs, but was applied to

every star detected by our pipelines.

This method allowed us to reduce the scatter in the ZPT values, and to determine

(by field and by filter) which was the cutoff radius from the minimal distortion center, at

which ZPT differences with respect to 2MASS rose above a maximum tolerance of 0.3

mag. Inside the area marked by this cutoff radius, the polynomial correction reduced the

ZPT differences significantly, and this also results in a decrease of the noise in the color

terms. The areas located beyond the cutoff circles, towards the east (left) corners of the

detector, have too large optical distortions, and so objects detected in those areas were

removed from our final catalogs

Figure 3-4 shows schematically the positions and extensions of every field observed,

as well as the cutoff radii of the ZPT correction for each filter. The K band circle, being

the most conservative, always defines the area of the field that was kept for analysis. This











of course, has an exception for those areas that have an overlap with the good quality

regions of another coincident field, in which case our catalog joining program selected

systematically the star from the good frame into the final catalog.






Rosette Complex
5.567111 1


5.12495




4.68279




4.24063




3.79847




3.35631
99.3568


98.9089 98.4611 98.0132
Right Acension (J2000)
FLAMINGOS GMC Survey


97.5654 97.1175


Figure 3-4: The extension of the areas of acceptable optical distortion are marked for
each field as circles with radii equal to the center of the maximum bin at
which the ZPT polynomial correction to the zero points (see text and fig-
ure 5) can be applied within the detector. This effect varies by field (size of
the acceptable area) and filter: the solid, dotted and dashed linestyle circles
represent the tolerance radii for J, H and K respectively.










In the various panels of Figures 3-5 to 3-9, we show, as an example, the effects of

the polynomial ZPT correction in the area 01 of our survey (which coincides approxi-

mately with the center of the Rosette Nebula)..

The first group of plots (Figures 3-5) shows the polynomial ZPT correction applied

to the photometric differences FLAMINGOS vs 2MASS as a function of radial position

from the minimal distortion center (3000,2400). As can be noticed, the scatter in the zero

point per magnitude is clearly corrected and the differences at large radii now converge

closer to zero and stay within a 0.1 mag range.


Figure 3-5: Example of the results of the ZPT polynomial correction of the zero point as
a function of pixel radial distance from the center of low optical distortion
(3000,2400) in the FLAMINGOS detector for region 1 of our Rosette survey
in filters J and H. The dots represent matches of the FLAMINGOS data with
2MASS sources in the ranges 11.0 to 15.0 and 11.0 to 16.0 H and J respec-
tively. The solid line represents a 6th order Legendre polynomial fit to the
median values of the scatter in bins of 300 pixels. The dashed lines indicate
levels of 0.1 mag of scatter.


-- -- -- -- -- -


..






















-0.5


1 .0 . . . .
0 1000 2000 3000
R (xc=3000, yc=2400)

2002 Jan 13 ROSETE 01 K see 1.00
1.0 --
0.000+-0.092
0.5


0.0


-0.5


1 .0 . __.. . .___ . ..___. .
0 1000 2000 3000
R (xc=3000, yc=2400)


Figure 3-6: Same as Figure 3-5 but for the K filter.


In the second group of plots (Figure 3-7) we can see how the values of the pho-

tometric differences with respect to the observed colors are also reduced intrinsically,

decreasing the overall uncertainty of our photometry.

The third set of plots (Figure 3-8) shows the net effect of the ZPT correction on

the color-magnitude and color-color diagrams. As it can be noticed, the color-magnitude

sequences for background and for members in the field get more confined and better

separated, which in the color-color space results in a reduction of the scatter around

the zero age and giant sequences. This, consequently, reduces the number of spurious

detections in the infrared excess region of the color-color diagram,especially in the area

located closer to the intersection of the T-tauri reddening band and the main sequence.

In the fourth set of plots of Figure 3-9 we show the scatter of these photometric

differences for field 01 before and after the correction is applied. In the left panels of the

















0.0 0.0


0.5 0.5

1.0 1.0
0.0 0.5 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 2.0 2.5 3.0 3.5

H < 15.00 & R < 3387 H < 15.00 & R < 3387
1.0 1.0


0.5 0.5
M* V


0.0 o0.0 .

0.5 -0.5

1 .0 . 1 .0 ._.._._.._._.._.
00 0.5 1.0 1.5 2.0 2.5 0.0 0 5 1.0 1.5 2 0 2.5 3.0 3.5
H K (2MASS) J H (2MASS)

Figure 3-7: Example of the results of the ZPT correction in the J -H and H K color
differences of FLAMINGOS vs 2MASS as a function of magnitude in area
1 of our Rosette survey. Top and bottom panels indicate color differences
before and after the polynomial correction. The color differences are also
reduced. The solid line indicates the zero level.


figure we plot the FLAMINGOS 2MASS differences as a function of magnitude, which

has larger scatter values towards fainter magnitudes. The mean value of the differences is

closer to zero after the correction. The right panels show how the correction results also

in a reduction of the net photometric scatter measured by the standard deviation of the

differences within consecutive magnitude bins. The median value of these deviations is

indicated with a dashed line.

3.5 Quality and Uniformity of the Survey

The overall uniformity in the quality of the photometry of our survey can be simply

assessed by comparing the mean values of the scatter in the FLAMINGOS vs.2MASS

differences after applying the positional correction, as described above. The median





















18 1i or ol
0.5 -
20 251 excess sources
0 .0 1 1 1 ,
0 1 2 3 4 -0.5 0.0 0.5 1.0 1.5 2.0
J K H K

counterfeit sigma PSF < 0.1 & R < 2705 counterfeit sigma PSF < 0.1 & R < 2705
3.0 -

2.5 -
14 / /r
St 2.0 /

S16 8 1.5 ..

S18 -1 .0 "

0.5
20 R < 3105 240 excess sources
. . . i . i . .0 0 1. 1 i . i .
0 1 2 3 4 0.5 0.0 0.5 1.0 1.5 2.0
J K corrected H K corrected


Figure 3-8: Example of the results of the zero point correction in the color-magnitude
and color-color spaces for the sources in area 1. The top and bottom panels
represent data and after before the polynomial correction is applied.


scatter values were compiled for every field and filter, and we use these numbers the

indicators of the net photometric quality of an observation. The medin scatter values for

each field are shown in Table A-2 of Appendix A.

The correction allowed us to reduce our internal photometric scatter in individual

bands by an average of 0.02 mag within the low optical distortion areas. The average

of these scatter values over the whole survey are 0.058 + 0.012, 0.064 + 0.018 and

0.056 0.014 in J, H and K respectively.

In addition to our completeness limits, which are statistical, there is another set of

limits which represent the sensitivity of the survey, i.e. the faintest magnitude at which

an object in our catalog can be consider to have good photometric quality within our

errors These sensitivity limits are also different for each filter, and we estimated them




















-^ - -
-0.5 0.05 -

0.00
10 12 14 16 18 10 12 14 16 18
JFLMN JFLMN
2002 Jan 13 ROSETTE 01 J see 1.00
1.0 0.25 -
-0 002+-0.071 fit lmnt llo 16o = 0.061
0.20
0.5
hs S a 0.15

^ --
0.0 0















sp1ec4tandvKe1ly.'7 ma.o
-0.5d a 0.05t0 estmatd c eten lits.I

0.00
10 12 14 16 18 10 12 14 16 18
JFLMN JFLMN


Figure 3-9: Example of the results of the ZPT correction in the values and scatter of
the FLAMINGOS vs 2MASS differences as a function of magnitude (J
band). Top and bottom panels show data before and after the correction re-
spectively. The numbers on the top of left panels are median and standard
deviations within the fit range. For the right panels, the dashed line and the
numbers on the top represents the median of the standard deviation per mag-
nitude bin. Solid and dotted lines mark the 0.0, 0.05 and 0.1 scatter levels.


simply as the points where the fiducial curve represented by the average values of the

FLAMINGOS photometric errors crosses the standard 107 limit defined as 0.109 mag

level in J, H and K. The result is shown in fig 3-10; our 107 crossing values are J=19.4,

H=18.4 and K=17.7 mag.

These values are also a good indication of the limits at which our photometric values

are consistent in quality, and are slightly higher than our estimated completeness limits. In

fact, it is possible that within certain areas of the fields,especially near the centers of the

low optical distortion areas, the completeness of the data is in fact higher than the average

values calculated from the artificial star experiments, but as stated in the section 3.7,













































Figure 3-10:


The solid fiduciall) line joins the points which mark the median value of the
PSF fitting photometric uncertainty in a given magnitude bin. Error bars in-
dicate the standard deviation in each bin. The horizontal dashed lines marks
the zero level and the sensitivity level, estimated at the standard 10-o limit
of 0.109 mag. Vertical dotted lines indicate the sensitivity limits, marked as
the bin at which the fiducial curve generated by the median values crosses
the 10-o line. The dash-dotted line represents the 3-0 level from the fiducial
line; any star in our catalog with errors higher than this levels were rejected
from the analysis.


this depends greatly on the variations of the extinction across the Molecular Cloud, and

logically, even the completeness limits can be compromised accordingly in certain spots.










3.6 Construction of Final Catalog

3.6.1 Intrinsic quality: 2MASS Addendum

Our photometric uncertainties also degrade relatively steeply at the bright end,

approximately below 11.0 magnitudes in J, H and K. This is because this coincides with

the level at which the counts per pixel in the detector reach values above 3 x 104 counts,

right at the limit of linearity and saturation.

In order to account for this effect, we rejected, for our analysis, every source in our

catalog with a magnitude brighter than H=1 1.0, as this filter seem to be affected slightly

worse by the saturation effect. To complete this end of the magnitude spectrum, we

added 2MASS sources to complete our catalogs within the range 5.0 to 11.0 mag in all

filters. The total addendum of 2MASS objects to our survey is 798 objects.

3.6.2 Survey Area Merging

A final catalog that includes the identifications, astrometry, photometry and basic

information about variability in overlap areas for all the sources accepted from our

survey, was created using a catalog merging code that put together catalogs for individual

fields and calculated the weighted average values for matches (duplicates) in overlap

areas. In the case where a match source was located in an area of high optical distortion,

null values were ignored. Most of the regions in the survey are adjacent, so that the

overlapping regions were usually small; in fact, the outermost 1 arcmin ribbon of the final

combined images was in most cases trimmed out of the final images, because Pinkack

only runs on the regions of the image that contained a combination of at least 50 percent

of the individual dithers. However, regions 4, G1 and G2 have large overlapping regions

within the RMC area and in those cases the selective matching came in handy.

After merging all of the individual catalogs, including the 798 2MASS sources for

the bright end, our final survey catalog contained a total of 153,266 objects.

For our analysis however, we restricted our study to those stars with an uncertainty

values below 3.0o above the median (see Figure 3-10). The total number of sources in










this selection is 146,868. Completeness limits at this level were not affected as 97.3%

of these 3-sigma rejections comprised magnitude ranges above J=18.50, H=18.25 and

K=17.50 mag respectively.

3.7 Intrinsic Detection Constraints

We consider important to mention that surveys with an emphasis on embedded

young populations, even with the aid of infrared detectors, cannot assure full detection

of all embedded cluster members. Most of young clusters are embedded in the remnants

of their original molecular gas cores, which are expected to carry massive amounts of

dust and dense gas, to the point that some objects (specially low mass protostars) will be

so highly extinguished that they will not be detected even in the near infrared, and thus

a fraction of the true number of cluster stars will be left off the counting. In addition,

cluster populations will always be mixed with a significant number of foreground

and background objects, whose contribution has to be estimated from the control

fields. Contributions of field stars are corrected by extinction effects, counted per

magnitude bin and subtracted from equivalent counts on-field. After these subtractions are

applied, final counting for the number of members in a cluster region are only statistical

estimates and cannot determine the membership of individual sources.

Another related factor that has to be taken into consideration is that low mass pre-

main sequence stars become fainter at older ages, with the consequence that at a certain

sensitivity limit they will just not be detectable (Carpenter, 2000). This sensitivity limit

depends on the age and mass distribution, as well as extinction, and the net effect is only

known accurately for a few regions.

In our figure 3-11, which is similar to figure 18 from (Carpenter, 2000), we try to

show the intrinsic limitations on mass and age detection that our RMC survey has. For

our completeness limit of K=17.25 we expect to detect stars well beyond the HBL if all

the stars were younger than 2.5 Myr and no extinction was present in the line of sight;

if stars are older, the range of observable stellar mass is reduced. In a typical molecular









DM=1 1.02, DM97 models


108








107


106








105
0.05


0.1 1.0


2.0


Log(Mass) [Mo]

Figure 3-11: Contours of equal K magnitude value as a function of stellar mass and
age at the estimated distance of the Rosette Molecular Cloud (1600 pc,
distance modulus=1 1.02). For the construction of these plots we used the
pre-main sequence evolution models of D'Antona & Mazzitelli (1997). The
solid lines represent iso-magnitude levels with no extinction, while dotted
lines represent the same levels with a Av=5.0 mag extinction. The contour
at K=17.25 coincides with the completeness limit of our FLAMINGOS
survey, and the dotted vertical line marks the 0.08 M. HBL limit.

cloud, typical extinction values range from a few to 50 magnitudes in visual wavelengths,

which would cause the completeness limits to compromise up to 5 magnitudes in K,

causing the range of detectable ages and masses to be even shorter in some areas.







47

Most of the known clusters in the Rosette are deeply or partially embedded in their

parental cores (see 5), which means that their ages could be possibly no older than 1-2

Myr (the most recent spectroscopically estimated age of NGC 2244 is about 2 Myr),

and thus, more than 80% of the embedded stars could be detected, depending on the

extinction level, but this effect is not uniform even at sub-cluster scales.















CHAPTER 4
NEAR-IR SURVEY: ANALYSIS AND RESULTS

4.1 Introduction

As discussed in Chapter 1, the study of embedded clusters in Giant Molecular

Clouds (GMCs) is of capital importance to understand the problem of star formation:

at the embedded stage, star clusters have not evolved significantly, and therefore their

densities and mass distributions are still close to the original fragmentation of their

forming cores. From observations of embedded clusters in nearby GMCs (d < 2

kpc), there is good evidence that a major fraction of the stars are formed in a cluster

environment, with rich clusters (100 members or more) clearly dominating over small

groups, as they contain more than 80% of the embedded stellar population in GMCs.

Unfortunately, available catalogs of embedded clusters are still incomplete at the

small cluster sizes. Systematic surveys of molecular clouds that are focused on the

detection of young clusters are usually limited to those regions with signposts of star

formation (like the presence of luminous IRAS sources), which lead to the discovery of

the richest clusters. Only in a very few cases has there been an specific search for low

density groups or distributed embedded populations, given that they are logically more

difficult to detect, especially if they are mostly composed of low mass stars, with large

spatial distributions and projected against a high background of reddened sources.

The RMC is a particularly active star formation region, with a large OB association,

NGC 2244, whose winds have generated an expanding HII region. One unsolved

problem is to determine if the shock front generated by this photodissociation bubble

was the principal trigger of the formation of the observed embedded clusters found by

(Phelps & Lada, 1997) in the adjacent, highly structured Rosette Molecular Cloud (RMC)

(Williams et al., 1995).










All of the PL97 clusters are associated with a luminous IRAS source and a massive

molecular clump, and until now, no other study has been able to determine the existence

of additional clusters. One of the main problems is the large distance to the Rosette:

the cloud is located at d = 1.6 kpc, roughly 4 times further away than nearby clouds

like Orion or Perseus and for this reason previous studies with limited photometric

sensitivity (depth) were unsuccessful at giving any new information on the distribution

of young populations in the Rosette. For example, the 2MASS survey is complete only

to K = 14.3 mag, which is not enough to detect low mass stars or deeply embedded high

mass stars. This makes difficult to detect highly embedded clusters, specially if their local

surface densities are low. Another difficulty is that the Rosette is located at a very low

galactic latitude (b = -1.8 2.0), which implies a very high density of field objects, and

thus, if a search for clusters was performed using monochromatic wavelength counts like

it was done for other clouds, then corrections for background contamination would be

large and difficult to apply.

Our survey is designed to address these problems by a) doing deep observations

of the region, capable of detecting stars close or below the HBL and thus improving the

detectability of previous studies, b) replacing the use of monotonic wavelength counts

by emulating the technique of Li et al., in which the detection of young populations

is done by photometric color selection, and c) using the method of Nearest Neighbors

(Casertano & Hut, 1985) to distinguish areas with surface densities intrinsically larger

than the overpopulated field.

Among the main goals of our survey are to study the characteristics of the known

clusters, to determine if there are more, and to study their distribution across the complex,

and in the context of NGC 2244. Because our observations are deep enough to detect

low mass stars we should be able to trace well the structure and extension of embedded

populations in the Rosette Complex, adding valuable information about the nature of

stellar nurseries in GMCs.










4.2 Analysis

4.2.1 The Nearest Neighbor Method

Single band monotonic detection of embedded clusters is done by subtracting

normalized control field counts (corrected by extinction) from counts towards the cloud

line of sight. This is expected to give the total number of expected members of a certain

region, and clusters are defined as regions with surface densities significantly higher than

the field. Unfortunately the method is biased because it will only be able to detect very

large or very dense clusters. Another problem is that the subtraction of field counts is

more difficult for more distant clouds, because members are fainter and are mixed with a

larger number of foreground and background stars.

The Rosette Complex is located 4 times further away than other star forming

complexes where systematic searches for embedded clusters have been performed, like

Orion or Perseus. Furthermore, the Rosette is located at a very low galactic latitude

(b = -2) and towards the anticenter of the Galaxy (1 = 210), which results in an

intrinsically high density of field sources located in the foreground and background

of the cloud. For example, one typical off-source field in the Rosette observed with

FLAMINGOS can have an average of 6x 103 sources, almost 5 times higher than a field

in Orion.

To detect embedded populations in the Rosette we applied a density selection

technique, the Nearest Neighbors Method (hereafter NNM), to distinguish populations

with surface densities above the uniform field levels. This method has already been

applied succesfuly to nearby clouds by Ferreira et al. (2005) and gives reasonable results

for a large distance region like the Rosette. However, we improved the method with

the use of a color selection to separate the youngest members in a field as near infrared

excess (IRX) objects, assuring the detection of embedded clusters by increasing their

probabilities of membership. We expect this combined approach to give a non-biased

detection of young stellar groups and to give insight into their nature at the same time.










Before describing our color selection, we will review briefly the terminology and

concepts from the NNM relevant to this paper. A more detailed description of the general

use of the method for embedded populations can be found at Ferreira et al. (2005).

The calculation of nearest neighbor densities to detect clusters in a crowded field

was proposed by Casertano & Hut (1985). They proposed to estimate the local surface

density of objects in a certain field from the individual, relative surface density of each

object. The generalized form of the individual density estimator is:


j-1
j =- (4.1)

where Dj is the distance of the star to its jth member. This estimator has only one

degree of freedom, the number j of neighbors used to calculate the local density. Accord-

ing to Casertano & Hut, the larger the value of j, the smaller the fluctuations in the local

density estimations due to local irregularities, which is very useful to determine extension

and structure of large systems. However, j also defines the minimum number of particles

in the smallest substructure to be considered, and so j should be small if structures looser

or smaller than typical clusters are to be detected. They showed that j = 6 was the mini-

mum number at which fluctuations could be acceptable for populations of the order of 30

to 1000 particles.

The NNM also allows the definition of a "density center" and a "density weighted

radius" which define cluster centers and cluster cores, respectively. The density center is

defined as the density weighted average of the star positions in a field:


Xd = X( (4.2)
-iPj(i)
And the density or core radius, Roreis defined as the density weighted average of the

distance of each star to X ,j:












Rcore = -i djlPj(i) (4.3)
Ci pj(i)
Using the NNM, clusters can be detected as regions where average individual pj

values are larger than those of a uniform control field, and ifj is small enough, groups

of the order of N = 101 stars should be detected without bias. If we follow the definition

of N 35 stars as a minimum number to represent a cluster (Lada & Lada, 2003;

Adams & Myers, 2001), then any groups with less members could be considered loose

enough as to account for a non-cluster (distributed) population.

For example, Ferreira et al. (2005) applied a j = 20 estimator to 2MASS catalogs of

nearby (d < 1 kpc) molecular clouds and confirmed locations and sizes of clusters with

radii as small as 0.3 pc and total number of members as low as 205 stars.

4.2.2 Detection of Embedded Populations

We calculated the 20th nearest neighbor densities for stars down to the completeness

limit in our final RMC catalog and the control fields, expecting to be able to distinguish

at least NGC 2244 and the seven clusters from PL971 as regions with densities unequiv-

ocally higher than the field. The result was that NGC 2244 is indeed, very well traced

as a high density region, as were the zone of clusters PL04 and PL05 in the core of the

cloud. Clusters PL01, PL03, and PL07 were also distinguishable but their apparent ex-

tensions were not more noteworthy than some groups of stars that rose above the 30 level

only because they coincided with "patches" of low extinction around the main molecular

gas emission. Finally, clusters PL02 and PL06, the smallest in the cloud, presented den-

sities below the 30 level because their small number of members (around 30) resulted in

lower than average 20th NN densities, and thus are not distinguishable among the noise

levels of the distribution.


1 We use from now on the nomenclature "PL01" to "PL07" to refer to these clusters.










Particularly, cluster PL06, which is associated with the B-type proto binary

AFGL961, contains large quantities of obscuring material near its center, which makes

difficult the detection of embedded faint members even in carefully constructed near-

infrared maps (e.g Aspin, 1998); any clusters like PL06 will be difficult to detect in a

large j scheme because their number of members will be intrinsically small.

Given these difficulties, we repeated the NN analysis using only infrared excess

(IRX) stars, to assure the detection of only the youngest component of the cloud embed-

ded population, to minimize the contribution of the field and to promote the detection

of deeply embedded, low surface density clusters. Infrared excess stars are not expected

to exist in the control fields as they located away from the star forming clouds, and so

the density of these objects should be always higher for embedded populations. Also,

as the fraction of IRX stars in an embedded cluster is larger if the cluster is younger and

therefore more embedded, so that high extinction regions could actually have higher IRX

overdensities, improving detection.

4.2.3 Infrared Excess Stars

In the J H vs. H K color-color diagram, IRX stars fall to the right of the

reddening band defined by the projection of the Classic T Tauri star (CTTS) locus

(Meyer et al., 1997), along the direction of the extinction vector (Cohen et al., 1981). The

fraction of stars with infrared excess emission in a cluster is known to decrease with time

as early stellar evolution leads to the destruction of disks by photospheric UV radiation,

but for a deeply embedded population with ages of 1 to 2 Myr like the one expected in

the RMC (the OB association NGC 2244 is estimated to have an age of 1.9 Myr and the

embedded clusters cannot be older), the circumstellar disk fraction will be significant

enough and IRX stars counted from JHK excess will trace well the presence of the most

recent episode of formation (see e.g Lada et al., 1996; Carpenter et al., 1997).

For our study, we define an IRX star as one with colors that place it 0.1 mag (5 times

the standard deviation of the H K uncertainty) to the right of the ZAMS and above









J H = 0.47(H K) + 0.46 which defines the lower limit of the Classic T Tauri Star
(CTTS) locus of Meyer et al. (see Figure 4-1).




2.5gion 1








2.0
I /







S,-K




0.0 0.5 1.0 1.5 2.0 2.5 1.0


Figure 4-1: The near-infrared color-color space. The thick dark solid lines represent the
loci of the zero age main sequence and the giant branch (Bessell & Brett,
1988). The thick colored line is the Classic T-Tauri locus (Meyer et al.,
1997), which in this diagram is extended to the right to large H K values
and above and below by its observational error. The other dashed lines rep-
resent extinction along the direction of the reddening vector indicated by the
arrow on the left. The shadow region indicates where, under our definitions,
stars with possible infrared excess emission fall. Stars falling in the regions
labeled as 1 and 2 colors are usually affected by spurious detections and high
photometric color scatter.










The first constraint avoids contamination from non-IRX stars with large H K

uncertainties located close to right edge of the Zero Age Main Sequence Reddening Band

(ZAMSRB). The second constraint helps us to avoid including objects that locate in the

region below the CTTS line. Sources fall in this region mainly due to high color scatter

(see section 4.2.4 below), however unresolved galaxies with large color dispersions

(Labb6 et al., 2003) or distant background galaxies with highly inclined reddening vectors

(Heraudeau et al., 1996) may have colors that fall in this region of the diagram. For

stellar sources, there are cases in which nebulosity can add a blue J H component to

background stars with a low extinction vector and push stars to the region, as shown by

Montecarlo simulations of Muench et al. (2001).

4.2.4 Magnitude Depth Restriction for IRX stars

The combined effects of variable seeing quality over the seasons, and variability

of the focus quality across the wide detector field of FLAMINGOS, resulted in a high

intrinsic dispersion of color values for our sample, which cannot be eliminated with

the zero point corrections or the uncertainty restrictions. In figure 4-2 we illustrate this

effect in a contour level color-color diagram all of the stars in our working catalog. In the

diagram, made with a nyquist box size of 0.1 mag, the lowest level shown represents the

mean color-color space surface density and each subsequent level represents a step of 1

standard deviation. There is a noticeable bloating in the dispersion of colors at both sides

of the ZAMSRB at the core of the diagram, near the regions of lowest extinction.

The color dispersion is larger for faint stars. The scatter in H K has a major incre-

ment at approximately K=15.75, where the field object density increases significantly. In

Figure 4-3 we show equivalent contour level color-color diagrams made with separate

samples for stars with K < 15.75 mag and for stars with 15.75 < K < 17.25 mag respec-

tively. The diagrams show that the scatter for the bright end bins is smaller than for the

faint end bins. Our calculations indicate that the standard deviation of the H K color

uncertainties is twice as large for the faint end bins (0.11 vs. 0.21 mag). This effect is










11.0

-0.51 ,
-0.5


0.0 0.5 1.0 1.5 2.0
H-K


Figure 4-2: Contour level color-color diagram for all stars in the FLAMINGOS RMC
survey within the restrictions described in section 2.1. The diagram was con-
structed using a sampling box of size 0.1 mag. The lowest level represents the
mean value of the object counts, at 825 dex 2. Subsequent levels represent
steps of 1 sigma (3550 dex2).


slightly worse for some fields which were observed under less favorable weather con-

ditions or at a higher than average airmass. There are also regions of the survey where

the scatter is smaller and the quality of the colors is kept to fainter limits (see section

4.3.2). The statistical cuts we present are mostly conservative and assure the uniformity of

our statistics across the entire survey area.


I I I I I I I I I I I I I ''''' 1 1 1 1 1 1


/

//


) I











a) 11

v di i.n tw a o- o r n /B
2 0 2 0 -




1.0 1.0

0 5 0 5 -

0 0- 0 0 .

-05 , I 0 5
-0.5 00 0.5 1.0 1 5 20 2.5 -0.5 00 0.5 1.0 1 5 20 2.5
H-K H-K

Figure 4-3: Contour level color-color diagrams for stars in the FLAMINGOS RMC
survey divided in two ample groups of brightness. Both diagrams were
constructed with a Nyquist box size of 0.1 mag. The diagram a) shows the
distribution of colors for 'bright' stars within 5.0 < K < 15.75 mag, and the
diagram b) is for 'faint' stars within 15.75 < K < 17.25 mag. The contour
levels start at the mean level (360 and 466 dex 2 for a) and b) respectively)
with subsequent steps of 1 sigma (1770 and 1910 dex 2 for a) and b) respec-
tively).


The high scatter in near infrared colors affects directly the calculation of the number

of IRX stars in the survey, which locate to the right of the ZAMSRB. We performed

Montecarlo experiments in which we simulated the colors of stars drawn from a model

population with an age of 1 Myr (D'Antona & Mazzitelli, 1997) located at the distance

of the Rosette and we added to our simulated stars, color errors and extinction similar

to those observed in the survey areas. We found that the resultant number of stars with

colors similar to those of IRX stars was 5 times larger for stars in the faint end.

Because we are basing our analysis in the detection of infrared excess sources, we

had to limit the primary aspect of our analysis, the identification of embedded clusters,

to those stars in the bright end of the sample to assure an IRX sample with a minimum

of contamination. However, these bright IRX sources are only helping us to trace the

location and rough extension of clusters, and although our color uncertainties are high,










individual K band magnitudes are still good within 0.1 mag, which allows us later to

include stars down to the HBL in the areas traced by the IRX sources and calculate

correct luminosity functions with a generous bin resolution of 0.25 mag.

Also, a photometric depth limit of K=15.75 means a sample still almost 1.5 mag-

nitudes deeper than 2MASS and is equivalent (for dwarf type stars) to a stellar mass

range of 0.09 to 0.18 M. for a population of 1 Myr embedded in a cloud with a typical

extinction of 0 to 10 visual magnitudes (D'Antona & Mazzitelli, 1997). Thus, we should

be able to count IRX populations slightly above the HBL for a typical young cluster.

4.2.5 Nearest Neighbor Analysis for Infrared Excess Stars

A preliminary exploration using false color image combos of the less populated

clusters, PL01, PL02 and PL06, indicated that the typical number of stars in a modest size

cluster that can be detected "by eye" within areas of bright nebulosity, usually coincident

with embedded cluster cores, could be of the order of 30 members. This is close to the

minimum number that defines an association of stars to be a cluster, and below that, any

groups could be considered a distributed population.

The expected JHK infrared excess fraction in an embedded cluster less than 3 Myr

old is 20 to 60 percent, so the number of IRX sources is much smaller than the number of

sources in the full catalogs, and cluster have to be identified with less neighbors. Because

of this, instead of j = 20 as in FL06, we selected a value of j = 10, which assures

the detection of clusters with less than 20 IRX members (60 to 100 total members if

the fraction of stars with circumstellar emission is 20 to 60%) but gives local surface

densities with 15% more accuracy than the minimum j = 6 described by Casertano & Hut

(1985). Also we are able to determine the existence of populations distributed in groups

almost three times less dense than the minimum expected for a cluster.

We detected a total of 116834 sources IRX sources under our definition in the

bright end sample. In Figure 4-4 we show their 10th Nearest Neighbor distributions of

























RX 10th Ne


Figure 4-4: Nearest Neighbor distributions for bright IRX stars. The top panel shows the
distribution of 10th neighbor distances. The bottom panel is the distribution
of 10th neighbor densities. In the top panel line A indicate the limit of dis-
tances shorter than 1.0 pc, while the dashed line B indicates the midpoint
value at 1.83 pc. At the bottom panel the equivalent limits in density space
are also indicated.

distances, Do1 and local surface densities p o. The mean value found for Do1 was 1.83 pc

which corresponds to a pio = 0.2 (')2. This limit is indicated in Figure 4-4.

For the control fields we found 3 sources that had IRX colors down to a maximum

brightness of K = 15.75. However, we added another 16 sources which fall to the right of

the reddening band below the CTTS line but would have IRX J H colors if reddened by










an average value ofAv = 5.0 mag, typical of the cloud regions. The mean 10th neighbor

density among these 19 sources was 0.18 (') 2 which compares relatively well with

the mean value 0.2 (') 2 of the distribution of distances in the survey areas, so that we

considered a round value of 0.2 (') 2 as a background limit, below which we cannot

assure that an IRX source has a density high enough to be distinguished from the field.

The minimum value of the Do1 distribution in the survey to be 0.145 pc, which

represents a density of 29.5 (') 2 and is a good estimation of the typical local surface

density in the central regions of RMC clusters. The midpoint between this minimum

distance and the mean is 0.987 (roughly 1 pc) which corresponds to a p o value of

approx. 0.6 (')2. This can be considered as a good estimate of the average embedded

cluster size in the RMC.

4.2.6 Identification of Clusters

In Figure 4-5 we show the location of IRX stars with the levels of density described

above. All of the known clusters seem to be traced well in this selection, and we confirm

that they are the main regions of star formation in the Rosette Complex.

In Figure 4-6 we present a contour level plot of the local surface densities calculated

with the NN method (j = 10). The contours were constructed using a Nyquist sampling

box of 90". We define a cluster as a region for which a closed contour at 0.2 (') 2 con-

tains at least 10 IRX sources. Using this definition we found, in addition to NGC 2244

and the seven PL97 clusters, 4 additional areas that arise as significant but have not been

studied before:

The first is a region at the "Core" of the cloud, to the east of cluster PL05 and south

of cluster PL04. This region, which we designate as RLE08, contains a large number of

highly reddened sources, differing from the clusters PL04 and PL05 which have a number

of sources already visible in DSS plates and can be considered partially emerged. RLE08

appears to be a more recent episode of formation in this ample zone of formation at the

























8 4.68(


O

U
o 4.32(





3.96(





3.60(




Figure 4-5:


99.0000 98.6800 98.3600 98.0400 97.7200 97.4000
Right Acension (J2000)
FLAMINGOS GMC Survey. NOAO/UF

The location of IRX stars in the Rosette survey with brightness K
15.75 mag. The "plus" symbols are IRX stars with 10th neighbor den-
sities higher than 0.2 (') 2, while black dots are stars with densities
below 0.2 (')2. We also indicate the expected position of the cen-
ter of NGC 2244. Contours indicate levels of CO emission in steps of
20K km s-1. The dotted line indicates the limits of the survey coverage.


center of the cloud, in which clusters PL04 and PL05 are the largest and most brilliant

clusters.

The second is a substantially large, highly reddened cluster located in the southeast-

ern edge of the cloud, designated as RLE09. Along with RLE08, these clusters are clear

examples of clusters which are located in regions with large extinction values and thus
















5.40000

R LE1



5.04000-

C C 2244 GC2

0V 0

8 4.68000 -
N PL02

.2
0I0

SI PLO












3.60000 1 I I I
99.0000 98.6800 98.3600 98.0400 97.7200 97.4000
Right Acens-on (J2000)
FLAMINGOS GMC Survey. NOAO/UF


Figure 4-6: Identification of clusters in the Rosette Complex. The contours indicate 10th
Nearest Neighbor densities and were constructed with a nyquist box size of
1.5 arcmin. Labels for individual clusters are explained in text. The dotted
thin lines indicate the 15.0 KO km s level of 12CO emission, which we use
to define the extension of the main molecular cloud regions.



have average surface densities comparable or lower than the field. However, they contain

a large number of red sources that are easily distinguishable in JHK false color composite

images and such a large number of IRX sources that they stand out clearly as embedded


clusters.










A third new cluster is located to the east of NGC 2244, in the region of the cloud

identified as NGC 2237, which is well known for its content of gas pillar structures

(Carlqvist et al., 1998), and thus we assign it to that name. The existence of this cluster

was also suggested by Li (2005) in their study of 2MASS data. NGC 2237 is distinguish-

able as a zone of high surface density in all star Nearest Neighbor counts. Other patches

in the fields that coincided with the Nebula areas also presented high densities but when

we applied the IRX color selection, NGC 2237 was the only one -besides NGC 2244-

that was confirmed to coincide with a cluster.

A fourth group designated as RLE10 is located North of NGC 2244, and although

it has a very low surface density compared to the rest of the clusters, it contains 13 IRX

sources, from which at least six have colors suggestive of very large infrared excess,

while the region has a small extinction value.

4.2.7 Properties of Clusters

We analyzed each cluster individually, isolating appropriate sub-regions that varied

roughly from 25 to 120 arcmin2 depending on the apparent extension of the clusters in the

maps. This way we were able to determine cluster structures with total areas of 6 to 60
-2
arcmin2.

For embedded clusters in the Molecular Cloud regions we calculated the extension

of a cluster as the area Ap inside the polygon defined by the 0.2 (') 2 contour in each

analysis box, and consider as potential members all of the sources (IRX and non IRX)

down to K=17.25 inside it. Equivalent radii, Reqcan be defined as /Ap/7u and can be

considered as standard estimates of the total extensions of clusters.

In Table 4-1 we present for each cluster, its center coordinates (as defined from

equation 2), core radii (equation 3) and equivalent radii. We also show the number of IRX

sources to K < 15.75 and the corresponding fraction it represents.

The IRX percentages in the nebula clusters NGC 2244 and NGC 2237 are intrin-

sically smaller than those in the embedded clusters, roughly 10 vs 18-76 percent. This










Table 4-1: Young Clusters Rosette Complex

Cluster RA DEC Rcore Requiv NIRX VNIX a IRXF b
ID center, J2000 [pc] K < 15.75
PL01 97.96 4.32 0.37 1.16 29-5 0.28
PL02 98.31 4.59 0.94 1.46 32-6 0.33
PL03 98.38 4.00 0.32 1.69 80-9 0.44
PL04 98.53 4.42 1.10 1.85 89-9 0.24
PL05 98.63 4.32 0.86 1.31 57-8 0.18
PL06 98.66 4.21 0.73 0.75 13-4 0.52
PL07 98.88 3.98 0.38 0.88 22-5 0.61
RLE08 98.56 4.32 0.99 1.30 49-7 0.33
RLE09 98.78 3.69 0.74 1.49 65-8 0.76
RLE10 97.78 5.27 1.19 1.15 15-4 0.32
NGC 2237 97.59 4.93 1.94 1.91 36-6 0.15
NGC 2244 97.95 4.94 1.56 2.30 62-8 0.12



aNumber of IRX stars with 10th Nearest Neighbor densities above 0.2 (')2.
bIRX fraction with respect to total number of stars with K < 15.75 inside 0.2 (')2 contour.


is due to an expected lower rate of disk survival in the presence of UV radiation from

numerous OB stars (Dolan & Mathieu, 1999), as well as disk evolution in older stars

which results in reduced circumstellar excess emission. For the clusters embedded in

the Molecular Cloud areas (PL01-PL07, RLE08 and RLE09) the large IRX fractions are

suggestive of ages of 1 to 1.5 Myr or younger (Haisch et al., 2001; Hillenbrand, 2006).

The core radii, Reore of the Rosette clusters (see equation 4.3) have a range of 0.3 to

2.0, with an average of 0.930.48 pc. The equivalent radii, Req range from 0.75 to 2.30,

with an average of 1.440.44 pc. The distributions of these size estimates are shown

in Figure 4-7. We also show in the figure the distribution of the Rcore /Req ratios, which

peak at 0.650.27 and have in two cases (clusters PL01 and PL03) values below 0.5. The

clusters NGC 2244 and NGC 2237 are extended and their core radii are too close in value

to their equivalent radii, so that we considered them equal. The distribution of core radii

and core to total ratios is consistent with the study of Ferreira et al. (2005), and suggests

that clusters, in most cases, have a tight center but with well extended edges.








65




4

_j 3

1-)


0
0 1 2 3
CORE



4






0 2 3
REQ





3
z2



0 0.5 1 1.5 2
RCORE/REQ



Figure 4-7: From top to bottom: distribution of core radii, equivalent radii and core to
equivalent radii ratios for the Rosette clusters


We constructed color-magnitude and color-color diagrams, which we show in Fig-

ures 4-8 to 4-19. In the K vs. H K color-magnitude diagrams we show the photometry

for all of the stars inside the corresponding 0.2 (') 2 contour, and mark separately those

with infrared excess. We include the ZAMS locus and a PMS evolution isochrone of 1

Myr, as well as extinction vectors corresponding to 3 times the mean value (Av) in the

cluster analysis box. Stars falling to the right of the isochrone are affected by extinction

towards the line of sight of the cluster, revealing their embedded nature. In the J H








66


vs. H K color-color diagrams, the same stars are located above the dwarf and giant

sequences along the reddening bands, with the IRX sources located to the right of the

MS reddening strip. Those objects located at or near the zero age sequences, which in

the color-magnitude diagram locate preferentially to the left of the isochrone, are most

probably foreground stars or evolved cloud members that coincide with the line of sight

of the clusters.


Ara (orcmin)
3 2 1 0 -1 -2 -
I '.' I '' '

''" ; ,' .'. .




C


.4
: o [ "' .


98.010 97.985 97.960 97.935
Right Acension ()


t"


.t


97.910


A-


Figure 4-8: a) K band image, b) control magnitude diagram, c) color color diagram and
d) Radial Density Profile for the area corresponding to cluster PL01. See text
for explanation.


IRX


- j












Ara (arrcmin)
2 0 -2






*.. .'- .
** ^ W, "* "





98.37 98.34 98 31 98.28 98 25
Right Acens. on
S* ., tl* 9

98.37 9834 9831 98.28 9825
Right Arension ()


Pt


Figure 4-9: Same as Figure 4-8, for cluster PL02.


The fourth plot in each panel are radial density distributions calculated with a

method of equivalent areas (see e.g. Muench et al. (2003)) for all stars down to K = 17.25

inside the analysis boxes and calculated from the cluster centers. In this plots we

indicate the core and equivalent radii calculated from the IRX Nearest Neighbor distri-

butions. With the exception of PL02, PL06 and RLE10, which are the clusters with the

lowest surface densities, the rest present well defined radial profiles, which unfortunately,

due to poor statistics cannot be fit successfully to standard King or Plummer cluster

models, but show well extended tails that in some cases (e.g clusters PL01, PL04, PL07,


%o


k

71,;r :
















Ara (orcmin)
4 2 0


2 4


-4
." C I
*
*-., : .. .. ;, .. ; dI




., ,


.''. .' i .
:



98,476 98 431 98.386 98.341 98.296
Right Acension (")


.'f


f .


o



'[a a ^^~~X s



o
I


Figure 4-10: Same as Figure 4-8, for cluster PL03.



RLE09) present well defined secondary bumps suggestive of structure. In the case of


the nebula clusters NGC 2244 and NGC 2237 the radial distribution profiles show a


slow decline that implies a negligible core peak, and might be suggestive of an extended


structure. However, the counts in each of the equivalent areas used to construct these


profiles are not corrected by background, and as the extinction is lower in the nebula,


these profiles might be showing the effect of field contamination.


- - . -, ,,-















Ara (orcmin)
4 2 0 -2 4
4.489 I




4 4 1
4 '



"- I- i ^ -1




437 -
U..'. ^".


-4-4
4339 .., -, .. I
98.621 98 584 98.546 98.509 98.471
R ght Acension (a )


Figure 4-11: Same as Figure 4-8, for cluster PL04.



4.3 The Fraction of Stars in Clusters

Under the assumption that the IRX are tracing the correct distribution of embedded

populations in the Rosette complex, we can use them to estimate the fraction of stars

that belong to clusters. We made our calculations inside the molecular cloud areas first,

to account for deeply embedded clusters only, and then for the whole survey area which

includes the emerged clusters located in the Nebula area.

The total number of IRX stars detected in the survey is 116934, out of which

63025 stars have NN densities larger than the mean, 0.2 (')2. A total of 43621 stars


't.4.1.
*i 'C-."6


* r^














Ara (orcmin)
4 2 0 2 4
4.388
.,- ..r .. ... ..
.. .. A, :'V .

4353-- .... -







S. t i2
4- --


98,702 98 667 98.632 98 597 98.562
Right Acension (a )


--------


Figure 4-12: Same as Figure 4-8, for cluster PL05.


are contained within the estimated areas of the 9 embedded clusters PL01 to RLE09,

which occupy a total of 242 sq. arcmin. The remaining 53923 stars have local surface

densities lower than the mean, and thus cannot be distinguished from the background

field.

The area of the molecular cloud covered by our survey was calculated as the one

contained inside integrated intensity contour levels higher than 15 K km s 1. This

area is equal to 2747 sq. arcmin, which means the clusters occupy roughly 9% of the

cloud. Inside the molecular cloud areas we counted 12411 stars with densities lower


,- ....... ,














3 2














98719 98
98.719 98


ara (orcmin)
0 1


2 3


-2
-







'1


4 .. ,;

694 98.669 98.644 98.619
Right Acension (')


Figure 4-13: Same as Figure 4-8, for cluster PL06.


than the mean, and 437 are stars with densities larger than the mean but not associated

with the cluster areas.

For background correction purposes, we use a scale factor equal to the ratio of the

non-cluster areas of the molecular cloud to the area of the control fields. Using this factor,

we expect to see a total of 9410 field IRX stars, which leaves a total of 738 IRX

sources in the cloud areas that are not associated with a cluster. From this, we estimate

that the fraction of stars in clusters in the Rosette Molecular Cloud is 865%.


~
..
.
r ;


~



i'


r-,


? ~t~~ X















Ara (orcmin)
3 2 1 0 1 -2 -3
4.037 .' '. ; I **









D
.* -






i




3937 .... I- -
98.932 98907 98.882 98.857 98.832
Right Acension (a)


Figure 4-14: Same as Figure 4-8, for cluster PL07.



If we repeat these estimates for the whole area of the FLAMINGOS RMC survey,

7308 sq. arcmin, we find that there is a total of 54923 sources associated with clusters

after including NGC 2244, NGC 2237 and RLE10 in the counts. The clusters occupy a

total area of 390 sq. arcmin, or 5.3% of the total survey areas. The number of IRX stars

with densities lower than the mean is 53923, and there are 819 IRX stars with high

densities but no association with a cluster. The scaled number of expected IRX field

stars from the control fields is 26117 in this case, which results in a total of 35919


it~a,

:P
..
"8 ~ did









73

a) b)

Ara (aricmin)
2 0 -2
438 .' *L .... I .o.

". ,.38 `L... -
.2




*2 9 98 5 9 5 0
4: *. ." .'.. 1 .











Right Acenslon () H-K
c) d)
F r 1 S a F e 8 o Radius from center (pc)8
0.0 0.5 2 .
10




























A i n I t ihce dd clter popu
s o c. T I a f



,* 4 : Q




0-

0 1 2 3 4 0 1 2 3
H-K Radius from center (arcmin)


Figure 4-15: Same as Figure 4-8, for cluster RLE08.


stars non-associated with clusters and a total fraction of 60+5% of stars associated with

clusters in the whole survey.


Another interesting result is that in the case of the embedded cluster population,

208+15 sources are contained in clusters PL04, PL05, PL06 and RLE08, at the "Central

Core" of the cloud, which corresponds to 48+3% of the total number of embedded

sources. This means that approximately half of the recent births in the RMC occurred at

the most dense region of the cloud, which coincides with the main zone of interaction

with the Nebula (Heyer et al., 2005). If the whole survey is considered, then the Central
with the Nebula (Heyer et al., 2005). If the whole survey is considered, then the Central












a)


Arc (orcmin)
4 2 0 -2 -4





S- ., ,-
.- ''.,: '. o ,












98.871 98 826 98.781 98.736 98.691
Right Acension (")
c*
98,871 98826 98.781 98.756 98.691
Right Acension (a)
c)


B",P XX.
: :
~~~
,
1


Figure 4-16: Same as fig 4-8, for group RLE09.


Core clusters plus the clusters in the Rosette Nebula, NGC 2244, NGC 2237 and


RLE10, account for 563% of the recent stellar formation in the Rosette Complex,


which suggests that the formation occurred in two main episodes which resulted in the

generation of the biggest clusters, and then, a number of secondary episodes resulted


in the smaller, remaining clusters which are distributed in the remaining areas of the


Complex.


x .. ... ..


0 1 2 3 4
H-K
d)

Radius from center (pc)
.0 0.5 1.0 1.5 2,0 2.t








\ l i ;.
*:/ \,

( i















2 Arc (a-cmin)
2 0 2
P -- --- *... *
.. '% ..


..
_, t .." ;, :'. ."" ""


*** *. .m
~.- *. .,'


:: 9,.^ *'*.;
U
4..,

_-q .'* @',- 4. '
. .


-I- I *. I 1 I I I *. .. I . .
97.81 97.79 97.77 97 75 97.73 97,71 9
Right Acension ()


















x IRx


ir.


-1 ..
~
,


i,


7.69


Figure 4-17: Same as Figure 4-8, for cluster RLE10.


4.3.1 Distribution of Sources with Respect to the Rosette Nebula

It is important to mention that with the possible exception of RLE10, NGC2237

and NGC2244, all of the clusters are associated with a massive molecular clump, which

confirms their deeply embedded stage. From this, it is clear that the Nebula and the

Molecular Cloud areas expose different episodes of formation. Also, high density IRX

stars in the RMC area are mostly confined to the limits of the cloud, while in the Nebula

area, the stars that trace cluster populations are already exposed out from the molecular


I I












a) b)


Ara (orcmin)
6 4 2 0 2 4 6


-. -



C. E


.. 14 -
2-

S. 16

.X -ux IRX ...

97.745 97 708 97,672 97.635 97 598 97,562 97.525 0 1 2 3 4
Right Acension () H K
c) d)

Radius from center (pc)
0.0 0.5 1.0 1.5 2.0 2,5 5,0









4A
j4-










S........ I ......... I ......... I ......... II........ I ......... I
0 1 2 3 4 1 2 3 4 5 6
H-K Radius from center (arcmin)

Figure 4-18: Same as Figure 4-8, for cluster NGC 2237.



gas, possibly showing a more evolved population which evacuated most of the molecular


material in the northern half of the complex.

We calculated the distribution of IRX sources with densities higher than the mean as

a function of the distance to the center of NGC 2244. To do this, we counted the number


of IRX sources inside the central parsec of NGC 2244 (11 sources) and then we counted


those outside this area in concentric annular wedge sectors with a constant width of 1.0

pc, but assuring that these sectors were always contained within the survey map areas. We


scaled and normalized the star counts in each wedge area with respect to the area of the









77

a) b)


Arac (aorcn)in)
6 4 2 0 2 4 6

| ., ,-, .- .-.. .- io:: o
-.'. S ,ooo
14






.. : -,.. .... i ---

'I16
.. .. 4



98.060 98005 97950 97895 97.840 0 1 2 3
Right Acension () H-K
c) d)
Radius from center (pc)
00 0.5 1.0 1,5 A2. 2.5 3,0






t 100-
Radus from center (pc) r in
OO 0.5 1,0 1,5 1 5 -,0











Figure 4-19: Same as Figure 4-8, for cluster NGC 2244.











first parsec circle. The result is shown in Figure 4-20, and we marked in the figure the

approximate locations of clusters and main cores of the molecular cloud.
*I I I 1I .



0 1 2 3 4 0 1 2 3 4 5 6











In the top panel of the figure, we see how the prominence of the Rosette Nebula

clusters indicate they are the largest stellar groups in the complex. The Molecular

Cloud "Rigde", where clusters PL01 and PL02 are located, and which is the part of

the molecular cloud that is in direct contact with the ionization front from the Nebula,

appears to be moderate in its star forming efficiency. The "Core" or central part of the

cloud, which contains most of its mass and which has been suggested as the main region














CLOUD CENTRAL CORE
PL03
PL06


Figure 4-20: Top panel: Distribution of IRX stars with NN densities higher than 0.2
(1)2 as a function of distance from the center of the Rosette Nebula
(NGC 2244). The counts are made in sectors of 1.0 pc in length and counts
in each sector have been scaled and normalized to the area and counts in
the central 1.0 pc circle in NGC 2244. Labels indicate the approximate
locations of clusters described in this paper, as well as the main 'regions'
of the complex. Bottom panel: equivalent distribution only for sources non
associated with cluster areas.

of interaction between the molecular and atomic hydrogen clouds (see Celnik, 1985;

Cox et al., 1990), and where the clusters PL04, PL05, RLE08 and PL06 are located,

seems to be carrying most of the cluster mode production, enhanced by the presence of










cluster PL03, which is however, located in a separated sub-cloud but at the same radial

distance. Clusters PL03, PL04, PL05, PL06 and RLE08 account fot 58% of the total

cluster population.

At the "Back Core" of the cloud, there are two clusters, PL07 and RLE09 which, al-

though smaller than those in the central core, still have significant extensions. Particularly

RLE09 has an extraordinary number of young sources despite its location well beyond the

interaction front of the Nebula. For these clusters, it is possible that a mechanism different

than triggering by interaction with the expanding HII region need to be proposed.

In the bottom panel of Figure 4-20 we repeated the counting but only considering

stars in the wedges that are located outside of the clusters. The sources not associated

with clusters were defined as those located at least two cluster radii away from each

cluster center. The scaled and normalized counts for these stars are of course much

smaller but it can be seen that there are two major zones of along the wedge where non-

cluster young sources accumulate: the first one is the area between the NGC 2244 and

the "Cloud Ridge", and the second one is the region in between the cloud "Central" and

"Back" cores.

4.3.2 A Case for a Distributed Population?

Noticing how there is a significant number of sources not associated with clusters

in the region between the central and back cores of the cloud, we used Field 09 of the

survey for a separate analysis. This field lies precisely to the south of the cloud "Central

Core" and north of the cloud "Back Core". The seeing and observing conditions for this

field were particularly good, and 93% of the original area of the field was kept after the

polynomial correction. The average scatter of colors down to K = 17.25 remains below an

acceptable 0.109 mag across the whole field, probably because the southeastern quadrant,

which for other fields presents high stellar profile distortions, overlaps in this case with

the good quality northwestern quadrant of the Gap 1 field. The weighted averaging