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MID-INFRARED HIGH RESOLUTION IMAGING OF HERBIG AE/BE STARS:
EXPLORING THE GEOMETRY OF CIRCUMSTELLAR DUST
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
@ 2007 Naibi Marifias
To my grandparents (mima, abuelo Raul, abuela Odilia y abuelo Paco)
I would like to acknowledge my advisor, Charles Telesco, for his love and enthusiasm for
astronomy. I will always be indebted to him for giving me the opportunity to work at the Gemini
Telescopes as part of the University of Florida mid-infrared instrumentation group. I also like to
thank him for allowing me work independently on this research while still providing guidance
when I needed it, and for trusting my decisions.
I am grateful to all the committee members for their support during the oral defense. I
specially want to thank Dr. Yasu Takano, for taking the time to carefully read and correct all the
typos and problems with the written dissertation.
I am very grateful to Chris Packham for being the best friend anyone could ask for, for
helping me move along in difficult times, and for reminding me that this was difficult, but
Many other friends also helped along the way: Lori read part of this work and corrected
my English, Joanna gave me invaluable advice on how to meet the deadlines during this last
semester, and Cynthia was always there to offer a hand when needed.
I also like to thank our secretary, Catherine Cassidy, for taking care of all the paperwork,
reminding me deadlines, and pointing out other things that I had to do in order to graduate.
I could never have finished this research without the support of my family: my mom who
came on weekends to take care of everything in the house, Ridel who also came every weekend
to take care of our son, Onelia who stayed with me this last semester so that I could write this
dissertation, and my aunt. I also have to give a very special thanks to my three-year-old son for
letting me know that even though he wanted to be with me, it was ok to stay with grandmas and
I am very thankful to the NASA Graduate Student Research Program (GSRP) and to the
South East Alliance for Graduate Education and the Professoriate (SEAGEP) for their financial
TABLE OF CONTENTS
ACKNOWLEDGMENT S .............. ...............4.....
LI ST OF T ABLE S ........._.._ ..... ._._ ...............9....
LIST OF FIGURES .............. ...............10....
AB S TRAC T ............._. .......... ..............._ 13...
1 INTRODUCTION ................. ...............15.......... ......
Star Form ation .............. ...............15....
Herbig Ae/Be Stars ................... .. ............ ... ...............19...
Geometry of Circumstellar Dust: Disk vs. Envelope .............. ...............20....
Observations ............ ...... ...............21....
Direct imaging ........._.. ........_. ...............21...
Spectroscopy .............. ...............21....
Interferometry ................. ...............23.................
P ol ariz ati on ................ ..... ............ ...............24.....
Modeling the Spectral Energy Distribution............... ..............2
Research Overview............... ...............29
2 MID-INFRARED INSTRUMENTATION, OBSERVATIONS AND DATA
REDUCTION .............. ...............38....
Infrared Sky .............. ... ........ ............3
Ground-Based Mid-infrared Astronomy .............. ...............39....
Atmosphere ................. ...............39.................
Gemini Observatory ..................... .. .............. .......4
Mid-infrared Instruments: Michelle and T-ReCS .............. ...............42....
Observing Technique: Chop-Nod............... ...............43
Output Data Files Structure ................. ...............44................
Output Image Spurious Structure .............. ...............45....
Data Reduction ................. ...............46...
Image Cross-Correlation .............. ...............47....
Flat Fielding ................. ...............47......_... ...
Sky Subtraction .............. .. ........ ... .........4
Standard Star Calibration and Airmass Correction .............. ...............48....
Photom etry .............. ...............49....
Color-C orrection ........._.. ........_. ...............49.
3 SOURCE SAMPLE AND DATA ANALYSIS .............. ...............59....
Herbig Ae/Be Sample .............. ...............59__ ......
Calculating Stellar Properties ............. ..... ._ ...............59...
Photospheric Emission .............. ...............60....
Line Emission ................. ...............62........... ....
Classification of sources ................. ...............62........... ....
Source Sizes ............... ... ...............63.......... ......
Point Spread Function .............. ...............63....
Moffat Function ................. ...............64.................
M easuring Sizes............... ........ .............6
Properties of the Circumstellar Dust. ........._... ......___ ...............66...
Dust Temperature .............. ...............67....
Warm Dust Mass ........._... ......___ ...............69....
Dust Size Distribution .............. ...............70....
State of Evolution of Dust Grains ........._... ....____......___ ..........7
Opacity ...................... ...............71
Passive Flared Disk Model .............. ...............72....
4 AB AURIGAE............... ...............84
Observations and Data Reduction .............. ...............85....
Source Size .............. ...............86....
Dust Properties ................. .. .......... ...............87......
Particle Temperature and Sizes .............. ...............88....
Dust Optical Depth and Disk Morphology ................. ...............90...............
M odel s .............. ...............9 1....
5 RESOLVING HERBIG AE/BE PROTOPLANETARY DISKS IN THE MID-
INFRARED: THE GROUP I SOURCES............... ...............104
Introducti on .................. ................ ...............104......
Observations and Data Reduction .............. ...............106....
Extended Emission .............. ...............107....
Notes on Individual Sources............... ...............108
HD36112 .............. ...............108....
HD 97048 .............. ...............109....
HD 100453 ............. ..... __ ...............111..
HD 135344 ................. ...............112......... .....
HD 139614 ................. ...............113......... .....
HD 169142 ............. ..... __ ...............114.
HD 179218 ............. .. ....__ ...............114
HD97048: An Alternative Scenario................ ...............11
Spherical Emission: Flaring Disk or Envelope? ................ ...............119..............
Geometry of the Systems: Disks or Haloes? ............ ...............120.....
6 RESOLVING HERBIG AE/BE PROTOPLANETARY DISKS IN THE MID-
INFRARED: THE GROUP II SOURCES .............. ...............138....
Introducti on ................. ...............138................
Notes on Individual Sources............... ...............139
Observations and Data Reduction .............. ...............144....
Extended Emission .............. ...............145....
Dust Properties ................. ...............148................
Discussion ................. ...............148................
7 CONCLUSIONS AND FUTURE WORK ................. ...............165........... ...
Disk Geometry in Herbig Ae/Be Stars .............. ...............165....
Future Work............... ...............166.
LIST OF REFERENCES ........._... ...... .__ ...............171..
BIOGRAPHICAL SKETCH ........._... ......._ ...............177...
LIST OF TABLES
3-1 Physical properties of program stars .........__ ....... ...............77.
4.1 Comparison of flux density measurements for AB Aurigae..........._.._.. ........._.._.......96
5-1 Log of observations for Group I sources .............. .....................125
5-2 Standard stars photometry............... ..............12
5-3 Derived flux measurements. ............. ...............125....
5-4 Quadratic subtracted sizes at 11.7 Clm for all Group I sources. .................. ...............126
5-5 Quadratic subtracted sizes at 18.3 Clm for all Group I sources ................... ...............12
6-1 Log of observations for Group II sources ....._..._._ ........_.. ........_ .........15
6-2 Standard stars flux densities............... ...............15
6-3 Derived flux densities .............. ...............155....
6-4 Derived dust color temperatures and masses for all sources included in this survey......155
LIST OF FIGURES
1-1 Stages of stellar and planetary system formation. ............. ...............30.....
1-2 H-R diagram showing the locus of different types of stars. ............. ......................1
1-3 Images of very young stars in the Taurus molecular cloud.. ..........__.. ....._.........32
1-4 Infrared spectra of the HAeBe stars Elias 1 and HD97048. ............. .....................3
1-5 Size-luminosity correlation derived from near-infrared interferometric observations..... .34
1-6 Typical spectral energy distribution (SED) of a HAeBe star. ............. .....................3
1-7 Schematic of the modified Chiang-Goldreich model .......... ................ .................3 6
1-8 Schematic of the flat disk plus halo model ................ ...............37........... .
2-1 ATRAN model of atmospheric transmission for Mauna Kea. ............. .....................5
2-2 Full width at half maximum measurements as a function of Universal Time ................... 54
2-3 Optical design for a Ritchney-Chretien telescope. ............. ...............55.....
2-4 T-ReCs optical design............... ...............55.
2-5 Filter transmission for T-ReCS (blue) and Michelle (magenta) ................... ...............5
2-6 Chop-Nod strategy ................. ...............57................
2-7 Spurious structure in the images ................ ...............58...............
3-1 H-R diagram for all the sources included in this survey ................. ............... ...._...78
3-2 Spectrum from stellar photosphere ................ ...............79........... ...
3-3 Plot of the ratio of near- to mid-infrared luminosity versus mid-infrared colors ........._....80
3-4 Normalized intensity profile of a Gaussian function and Moffat functions with
different parameters 0 .............. ...............8 1...
3-5 Representation of the Chiang and Goldreich flaring disk ................. ......._.. ........._82
3-6 The two solutions for the modified Chiang and Goldreich disk model.................... .........83
4-1 Near-infrared scattered light of AB Aurigae taken with STIS at the Hubble Space
Telescope .............. ...............97....
4-2 Near-infrared scattered light image of AB Aurigae showing spiral structure ..................98
4-3 AB Auriga false color images ................. ...............99...............
4-4. Normalized contour level of AB Aur, PSF star and PSF subtracted emission ........._.....100
4-5 11.7 Clm data (top panel) and 18.1 Clm (lower panel) FWHM for AB Aur (solid
circles) and PSF star (open triangles) .............. ...............101....
4-6 Ratio of the thermal emission at 1 1.7 and 18.1 Clm shown for different values of 10 .....102
4-7 Radial distribution of the dust temperature from the flaring disk model ................... ......103
5-1 Clearly resolved sources in the sample ...._.. ...._._._._ .........__. ...........2
5-2 FWHM measurements for HD36112, HD139614, and HD 179218 .............. .... ........._..128
5-3 Residuals after subtraction of the normalized PSF in the Si-5 filter. .........._... ..............129
5-4 Group I sources and PSF stars in the Si-5 filter............... ...............130
5-5 Group I sources and PSF stars in the Qa filter ................. ...............131...........
5-6 Modeled and observed emission from a flared disk with an inclination of ~ 45
degrees. ............. ...............132....
5-7 HD97048 cross-convolved Gemini images .....__. ............... ........._.._.......13
5-8 Location of the Herbig star HD97048 on the southeastern edge of the Chameleon I
dark cloud............... ...............133.
5-9 IRAS images of the Ced 111, reflection nebula associated with HD97048 ..................134
5-10 H band polarization map of Chal IRN ................. ...............135.............
5-11 1.3 mm observations of HD97048 showing the offset of mm emission relative to the
optical and near-infrared source position. ....._ .....___ ........__ ............3
5-12 Offset of the Si-5 and Qa peak of emission relative to the optical position of
HD97048 (left panel) and PSF star (right panel) ................. ...............136.............
5-13 Ratio of the thermal emission at 1 1.7 and 18.3 Clm for different values of lo ................ 137
6-1 FWHM sizes versus UT time. ..........._ ..... ..__ ...............156.
6-2 Residual emission at 11.7 Clm. ............. ...............157....
6-3 Photometry from the literature and from this study for the companion source to
HD 144668 (Rossiter 3 93 0)........... ......__ ...............158
6-4 H-R Diagram showing the position of HD 144668 (largest circle) and Rossiter 3 930O
(sm all est circl e) ........._...... ...............159.__..........
6-5 Plot of the ratio of near- to mid-infrared luminosity versus mid-infrared colors ............160
6-6 Measured sizes at the FWHM level in the Si-5 and Qa images ................. ................. 161
6-7 Ratio of warm to cold dust disk mass for all the sources with available cold dust
mass values in the literature ................. ...............162........... ...
6-8 Correlation of sizes and fraction of extended to total emission as a function of age. .....162
6-9 Group I and Group II sources in the H-R diagram .............. ...............163....
6-10 Color temperatures for all the sources included in this survey as a function of stellar
ages .............. ...............164....
7-1 Sketch of the evolution of protoplanetary disks in HAeBe stars ................ ................. 169
7-2 Spatial resolution of Spitzer (blue line) from near- to far-infrared wavelengths and
the IRAS (red lines) spatial resolution ................. ...............170..............
Abstract of Dissertation Presented to the Graduate School
of the University of Florida in Partial Fulfillment of the
Requirements for the Degree of Doctor of Philosophy
MID-INFRARED HIGH RESOLUTION IMAGING OF HERBIG AE/BE STARS:
EXPLORING THE GEOMETRY OF CIRCUMSTELLAR DUST
Chair: Charles M. Telesco
Herbig Ae/Be (HAeBe) stars are emission line pre-main-sequence stars of intermediate
mass. Circumstellar dust is well established as the origin of the large infrared excesses
characteristic of these stars; however, the geometry of this dust has remained controversial for a
long time. The sources have been divided into two groups depending on the infrared excess:
Group I sources have strong near- to far-infrared excesses, and Group II sources, only have
strong near-infrared excess. Previous studies associate Group II sources with optically thick flat
disks and Group I sources either with optically thick flat disks surrounded by haloes or flaring
disks with an inner hole at the dust sublimation radius and a puffed up inner wall. These two
models predict the same overall flux at all wavelengths. An evolutionary scenario has been
proposed in which grain growth, dispersion, and settling of larger grains towards the mid-plane
change the geometry of circumstellar dust from flaring disks in Group I sources, to flat cold
disks in Group II sources in about 1 Myr.
For this study, I imaged a sample of 20 HAeBe stars (8 Group I and 12 Group II) in the
mid-infrared from the 8-meter Gemini telescopes. I resolve extended emission in all Group I
sources with sizes tens to hundreds of AU. Most of the resolved emission is spherical, which
could be interpreted either by a halo or the surface of an almost face-on flaring disk; however,
the distribution of dust grain sizes that I derived from these observations, with larger grains at
smaller radii is inconsistent with halo models. Only three of the Group II sources are extended
and two of them show emission consistent with highly inclined disks. I find no correlation
between stellar ages and geometry of the circumstellar material supporting an evolution from
Group I sources to Group II sources, on the contrary, Group II sources are younger than Group I
sources and all resolved sources presented here are older than 1 Myr.
During the last two decades and motivated by a desire to understand the origin of our own
solar system, a significant effort has been placed into the study of pre-main-sequence (PMS)
stars. These studies combined with the advances in stellar evolution modeling, have allowed us
to integrate the observational properties of young stars at different wavelengths and create a
coherent theory of the physical processes involved in the formation and evolution of stars and
planetary systems. In this research I discuss the properties and geometry of dust around young
stars of intermediate mass. These circumstellar structures of gas and dust are the remnants of the
stellar formation process and the sites of planetary systems formation.
We can divide the stellar formation process into four stages. During the first stage of star
formation, when the balance between pressure and self-gravitation is lost, dense cores form in a
massive, cold, nearly static molecular cloud. The cloud starts to collapse, fragmenting into
smaller (~ few solar masses) regions. The resulting fragments continue to collapse on their own
in a quasi-isothermal mode because the densities are too low to prevent cooling. Most of the
energy released during the gravitational collapse is radiated away and absorbed by the optically
thick envelope, keeping the interstellar gas in an envelope surrounding the core almost
isothermal. The spectral energy distribution of the system at this stage is the signature of the
collapsing envelope and peaks in the mid- to far-infrared wavelengths (Figure 1-1, top row).
When the density and temperature of the core become coupled and the collapse becomes
adiabatic, fragmentation stops and the different cores evolve into protostars. The density at the
core determines the rate of matter infall onto the core. The timescale for this process, t, is given
where G is the gravitational constant and pc is the density of the core. Eventually, when the
pressure at the central core becomes significant, the core reaches quasi-static equilibrium and
becomes a protostar on the Kelvin-Holmholtz timescale given by
where M~c, Rc, and Lc are the mass, radius and luminosity of the core respectively. During the
second stage, gas and dust closer to the core accrete faster onto the protostar than the mass at
larger radii, and pressure gradients develop. Matter originating far from the rotation axis (defined
by the average angular momentum of the circumstellar material) has too large an angular
momentum to fall onto the protostar and begins to settle into a disk. The high accretion rates
power bipolar outflows, which are aligned with the rotation axis of the system, and material near
the poles in the envelope starts to clear. At this point the spectral energy distribution of the
system has two components. It is still dominated by the circumstellar material peaking in the
mid-infrared; however, the protostar is beginning to emerge at shorter wavelengths (Figure 1-1,
In the third stage, accretion rates decrease, the envelope is mostly lost, and circumstellar
gas and dust are confined to a massive optically thick disk. This is thought to be the birthplace of
planets; therefore, the disk is called a protoplanetary disk. In the center of the system, the
protostar continues to contract until the core temperature is high enough to ignite nuclear
reactions. At this point the release of nuclear energy stops the collapse on the core and the
protostar becomes a star, reaching the zero age main sequence (ZAMS). The spectral energy
distribution of the system is now dominated by the optically visible star and peaks in the UV;
however, there is still a significant contribution of emission from dust at longer wavelengths
(Figure 1-1, third row).
During the last stage, the physical properties of the star remain mostly unchanged while the
star burns the hydrogen available in its core. The star becomes a main sequence star. In the
circumstellar disk, solid particles in orbit around the star settle into a layer in the mid-plane of
the disk and grow by collisions. The dust becomes optically thin to radiation from the star and
very little, if any, gas remains in the disk. At this point, the protoplanetary disk is called a debris
disk since most of the dust in the system has been reprocessed. Main sequence stars surrounded
by debris disks are sometimes called Vega-type stars. In the spectral energy distribution of these
sources, there is very little excess emission in the mid-infrared due to a remnant dust disk (Figure
1-1, bottom row). Growth continues through settling of small dust particles onto larger particles
and the gravitational pull of neighboring material. Some planetesimals become large enough to
accrete gas and form giant planets.
One would expect to be able to extend this theory of star formation to intermediate (~ 2 to
8 solar masses) and high mass (> 8 solar masses) stars; however, we have seen that the
timescales controlling these processes depend mostly on the mass of the original protostar.
Therefore, the timescales for the evolution of the central star and its surrounding material are
different. If tK > ty the protostar will still be contracting when the high accretion from the
envelope stops. This happens for stars of mass lower than 8 solar masses. However, if tff> tKH,
the core will evolve faster than the free-falling envelope and the star will reach the ZAMS while
still embedded in an accreting envelope.
There are also two different mechanisms for burning hydrogen in stars depending on the
initial protostellar mass. The proton-proton (PP) cycle uses only hydrogen to fuse helium. Since
hydrogen has only one proton, the temperature required to start this process is only about 10
million K, and the energy generated in this process is proportional to T4. This cycle predominates
in low mass stars (M < 1.5 solar masses). More massive stars need more energy to balance self-
gravitation and stop the collapse. If a more massive protostar has carbon, nitrogen and oxygen
present in its core, when it reaches temperatures of ~ 16 million K, it initiates the carbon-
nitrogen-oxygen (CNO) cycle. The energy released in this process is more sensitive to
temperature (E ~ T20) and can balance the contraction of the star. The different internal
temperatures of low and high mass stars also cause different mechanisms of energy transfer in
their cores. In more massive stars temperature changes rapidly with distance so that
act,,al > -adnbatic and the fluid becomes unstable and boils transferring energy by
convection. Cores of lower mass stars remain adiabatic and energy transfer occurs by radiative
On this basis, three types of PMS stars can be defined: (1) low mass (M < 2 solar masses)
stars, also called T-Tauri stars, have radiative cores and are optically visible during their PMS
phase; (2) intermediate mass stars (2 to 8 solar masses), also called Herbig Ae/Be (HAeBe) stars,
have convective cores and are visible during part of their PMS phase; (3) massive (M > 8 solar
masses) stars have convective cores and remain invisible in the optical during the whole PMS
phase. The intermediate mass stars represent an interesting transition between the well-studied T-
Tauri stars and their high mass counterpart and are the subj ect of this research.
Herbig Ae/Be Stars
In 1960, Herbig discovered a class of emission line stars of spectral types A and B
associated with nebulosity and claimed them to be the high mass counterpart of T-Tauri stars.
All the stars were located in obscure regions (star forming regions), had spectra with emission
lines, and were illuminating bright reflection nebulae (rej ecting nebulous objects).
Subsequent studies extended the membership criteria to include stars of spectral type F8 or
earlier and provided better continuity with T-Tauri stars (Finkenzeller and Mundt, 1984). They
also proved the PMS nature of these stars. Comparison of the stellar radial velocity of 27
HAeBe stars from photospheric Balmer lines to radial velocities of molecular clouds from
molecular lines showed that there was no systematic stellar motion relative to the clouds, and the
stars were still associated with the star forming regions (Finkezaller and Jankovics, 1984). Palla
and Stahler (1993) revised early evolutionary models and the role of deuterium burning
disproving the idea that stars heavier than 3 solar masses would not be observable during their
PMS phase, and showed the observational birthline (where stars first appear as optically visible
sources) for intermediate mass stars in the Hertzsprung-Russel (HR) diagram' (Figure 1-2).
Calculation of effective temperatures, surface gravities and stellar luminosities (Strom et al.
1972; Cohen and Kuhi 1979; van den Ancker et al. 1998) allowed the positioning of HAeBe
stars in the H-R diagram where they appeared above the zero-age main sequence (ZAMS) in
evolutionary tracks still moving toward the main sequence (MS). The strong ultraviolet and
infrared excess discovered by Mendoza (1966) could not be explained by free-free, free-bound,
and bound-bound emission in hot gas, but required dust thermal emission to account for the
SThe HR diagram is used to identify the evolutionary status of a star. It plots the effective temperature (Teff) of the
star on the x-axis versus the luminosity of the star (L) in the y-axis on a logarithmic scale. During its lifetime a star
changes position in the diagram following evolutionary tracks dictated by the mass of the original protostar.
larger excesses beyond 1 Clm. The detection of the dust related silicate feature in the 10 Clm band
of some HAeBe stars offered further evidence of the presence of circumstellar (CS) dust clouds
(Cohen 1980, Berrilli et al. 1992).
During the 1980s, the number of candidate HAeBe stars increased significantly when
Infrared Astronomical Satellite (IRAS) found a large number of stars resembling HAeBe, but
located in isolated regions. A catalogue of HAeBe stars by The et al. (1994) included 287
sources with the following membership criteria:
Spectral type B to F8
Near or far infrared excess due to circumstellar material
With or without:
Association with obscure star forming region
Irregular photometric variations
Variable or high degree of linear polarization
Geometry of Circumstellar Dust: Disk vs. Envelope
The presence of cool circumstellar material in HAeBe stars is well established as the origin
of their characteristic large infrared and sub-millimeter excesses; however, the spatial
distribution of this material remains controversial. The intrinsic complexity expected in these
more massive environments makes multiple interpretations of the same observational data
possible. In addition, there seems to be an equal amount of observational evidence favoring
spherical and axis-symmetric geometries.
Imaging of CS material at different wavelengths directly reveals the geometry of the
systems; however, these observations are limited by the spatial resolution of the telescopes used.
In the sub-mm, imaging of HAeBe star in continuum and molecular line emission (Mannings and
Sargent 1997, 2000, Corder et al. 2005, Pietu et al. 2005) showed gas emission regions extending
hundreds of AU in size. Masses of the CS material from these observations were in the range of
0.01 0.05 solar masses, and because there is low optical extinction toward the stars in all these
systems, this provided strong evidence for the existence of optically thick disks where most of
the cold mass could reside without covering the central star.
Optical and near-infrared scattered light images from Hubble Space Telescope (HST)
revealed large (hundreds of AU) disks in some sources, sometimes immersed in more extended
nebulosities, Figure 1-3. Structures (gaps, spirals) that signal the presence of large bodies
clearing the dust or gravitational instabilities in the disks were also visible (Grady et al. 1999,
2001). Mid-infrared imaging from large ground-based telescope has also resolved extended
emission in a few sources (Fisher et al. 2000, Polomski et al. 2002, Jayawardhana et al. 2001).
A study of emission lines of HAeBe stars by Hamann and Persson (1992) found a
correlation between excess infrared luminosity and Ca II line strength. The correlation could be
explained if the lines originate from material accreting onto the star or from fully ionized matter
very close to the stars evaporating due to intense thermal and radiative pressure. Lorenzetti et al.
(1994) explained the correlation between Hot and L band (3.5 Clm) excess in a similar fashion,
either resulting from winds powered by accretion or from very small excited superheated grains
emitting in the near-infrared as they cool down to equilibrium temperatures.
The presence of only blue-shifted wings in some forbidden line emission in HAeBe stars
has also being attributed to a disk obscuring the red-shifted flow. The study of [SII] by Corcoran
and Ray (1994) revealed some optical veiling (presence of non-stellar continuum) in the more
embedded stars supporting accretion models. The small percentage of detection was attributed to
the increased brightness of the stellar continuum of HAeBe stars in comparison to T-Tauri stars,
which may reduce the chance of revealing forbidden line emission. However, a similar study of
[OI] forbidden line emission showed that most of the surveyed sources had unshifted symmetric
line profiles implying an absence of CS disks at the scale mapped by the [OI] forbidden line
emission (Bohm and Catala 1994). More recently, Acke, van den Ancker and Dullemond (2005)
studied the origin of this more spherical [OI] forbidden line emission and found a strong
correlation of line strength with 60 Clm continuum excess emission. They interpreted their
findings in the context of a flared disk model. In this scenario UV radiation from the star can
penetrate more deeply into the flared disk, photodissociation of OH molecules in the disk and
non-thermal excitation of oxygen in the atmosphere of the disk were able to reproduce the
Near-infrared spectroscopy from Infrared Space Observatory (ISO) showed that while the
near-infrared excess is similar for all HAeBe stars, there are large differences in the mid- and far-
infrared excesses prompting the separation of these sources into two subgroups, with Group I
sources having strong near- and far- infrared excess, and Group II, only strong near-infrared
excess (Meeus et al. 2001). They also showed that the infrared spectrum of these sources is
dominated by amorphous and crystalline silicates, and in some cases by polycyclic aromatic
hydrocarbon (PAH) emission, Figure 1-4. Silicates have been observed in many young stars; the
general idea is that amorphous silicates from the interstellar medium (ISM) are crystallized
within CS disks. Van Boekel et al. 2003 showed a relationship between the silicate emission
strength and shape of these systems. The change in shape from a narrow peak to a broader
feature was explained by an increased in grain sizes and larger amount of crystalline silicates.
PAH emission lines are only seen in Meeus group I sources (those with stronger excesses) and
are thought to originate in optically thin dust around the star (Acke and van den Ancker 2004),
either as an optically thin halo or the surface of a flaring disk.
Recently, mid-infrared spectroscopy with high spatial resolution has revealed radial
variations of dust composition with crystalline silicates dominating the regions closer to the star
and a mixture of crystalline and amorphous silicates at larger distances (van Boekel et al. 2004).
Interferometry combines the light waves captured by multiple telescopes to achieve higher
spatial resolution. Near-infrared interferometry is used to probe the hot circumstellar dust very
close to the star. Millan-Gabet (2001, 1999) resolved the inner regions of a group of HAeBe
systems with typical source sizes of 0.5-0.9 AU. Their observations were consistent with
symmetrical brightness distribution in all resolved sources and since we expect different
inclination angles in disk systems, they ruled out the presence of disks. They also found that the
near-infrared excess was similar in all sources independent of the sizes of the emitting regions.
More recently, interferometric observations with larger baseline are showing evidence of
inclination in these systems (Eisner et al. 2004, 2003) favoring the presence of disks. Eisner et al.
(2004) and Monnier et al. (2005) also found a strong correlation between the derived sizes and
the stellar and accretion luminosity of the sources, Figure 1-5. The correlation can be explained
if the source size is set by the dust sublimation temperature, which depends on the stellar and
accretion luminosity and on the properties of the dust (Tuthill 2001, Natta 2001, Muzerolle
2003). The observations can be explained if the near-infrared emission originates on curved inner
walls at the dust-destruction radius (Monnier et al. 2006). Mid-infrared interferometry has also
resolved extended emission in these sources with sizes for the extended warm dust of a few 10
AU at most (Leinert et al. 2004).
Polarimetry provides additional information about the geometry of these systems because
CS material causes polarization in the light from the system due to scattering from dust.
Resolved extended axisymmetric geometries show spatial changes in the polarization vector,
unresolved axisymmetric geometries viewed with some inclination angle will have some degree
of polarization, while unresolved disk-like structures viewed face-on and spherical distributions
produce polarization vectors that cancel each other out resulting in a zero net polarization for the
system. Optical linear polarization (Maheswar et al. 2002) and Ha spectropolarimetry
observations (Vink et al. 2005, 2002) of HAeBe stars provided evidence for the presence of
disk-like CS structures. Change in linear polarization was found in a group of sources with
earlier types having a broad polarization consistent with a small scale (~.07 AU) flattened
structure, and later types (A2-F) showing a line polarization characteristic of a compact Ha
emission region being polarized by a rotating disk-like medium. The majority of the sources
with known outflows or companion stars had polarization vectors close to perpendicular or
parallel to the outflow or binary position angle.
Modeling the Spectral Energy Distribution
Modeling the observed spectral energy distribution (SED) of the sources is one way to gain
insight as to the geometry of the CS material. Assuming radiative equilibrium, the flux from the
central star absorbed by a grain at a distance r equals the power radiated by the grain, which
allow us to calculate the radial temperature distribution, T(r), of the CS material. With some
knowledge about the density distribution and physical properties of the grains, the SED of the CS
material can be computed for different geometries.
Hillenbrand et al. (1992) modeled the SED of 47 HAeBe stars and proposed an
evolutionary sequence based on the slope of the SED in the infrared: group II sources had flat or
rising spectra and were modeled by a star plus accretion disk surrounded by a spherical
envelope; group I sources had spectral slope AlF, ac Ai and could be fitted by using only a star
plus a flat accretion disk; and group III sources had very little infrared excess emission that could
be explained by free-free emission from an ionized region around the star, and did not required
dust thermal dust emission to be present. The proposed evolutionary sequence in this
classification was group II-group I-group III (Natta et al. 1998). In order to fit a 3 Clm peak
observed in the spectra of early type HAeBe stars with flat disks, Figure 1-6, the model required
high accretion rates (6 x 10-7 to 8 x 10-' solar masses per year) and inner holes of 5 to 20 stellar
radii in the disks so that the observed excess near-infrared luminosity in these systems originated
from accretion of circumstellar dust onto the central star.
The main problem with this model arose from the high accretion rates required for the
models to keep the region inside the inner hole optically thin. Clearing of grain dust around a star
is expected inside the dust sublimation radius, the distance at which grains are heated to ~ 1500
K and annealed; however, this inner region will remain optically thick at the high accretion rates
needed to fit the models (Hartman, Kenyon, and Calvet 1993). For an accretion disk, the radial
temperature distribution is set by the luminosity of the central star and the accretion luminosity.
The disk temperature increases with decreasing radius and since the opacity of the material
increases with increasing temperature, there should be an optically thick inner region, even
inside the dust sublimation radius. In addition, these high accretion rates were incompatible with
observations. To eliminate this problem, Hartman, Kenyon and Calvet (1993) proposed three
different scenarios as possible sources for the 3 Clm emission peak: 1) thick infalling envelopes
with empty cavities to provide a clear line of sight to the hot central region producing the
emission (model with very limiting parameters for all sources); 2) presence of companion plus
envelope (which would require all HAeBe stars to be binaries with clear line of sight to the
central HAeBe star and high extinction towards the companion to explain the high frequency of
the emission peak); and 3) dusty optically thin envelopes with a population of very small
transiently heated dust grains. This last alternative model explained the 3 Clm emission in early
type HAeBe stars and its absence in later spectral types since F-type stars do not produce enough
UV photons to heat large populations of very small grains to high temperatures.
Passive disks (flat and flared) were also invoked to explain the observed SED (Lada and
Adams 1992). Even though symmetric dust and gas envelopes were not the only possibility, they
were still required to explain the infrared component of the spectra in some HAeBe stars (Berrilli
et al. 1992, Pezzuto et al. 1997, Di Francesco 1994, Evans and Di Francesco 1995) and a
combination of geometries at different spatial scales seemed to be the only approach (Natta et al.
1993, Malfait et al. 1998).
Until this point, all the models assumed radiative equilibrium, and power-law
approximations for the dust density distribution. Chiang and Goldreich (1997) were the first to
self-consistently calculate the disk structure by considering hydrostatic and thermal equilibrium
in a passively irradiated disk. The gravitational force acting on a particle around a star has two
components: the horizontal component forces the particle to move in keplerian orbits around the
central star, while the vertical component pulls the particle towards the disk mid-plane. The net
result of gravity on circumstellar material is a flat keplerian disk. However, protoplanetary disks
also have large amounts of gas, and gas pressure can counter the vertical component of the
gravitational force resulting in a disk scale height that depends on the balance between gas
pressure and gravity. The Chiang-Goldreich disk model results in a flaring disk with an inner
optically thick layer where most of the disk mass is concentrated, and an optically thin surface
layer. The disk flares and intercepts more stellar radiation than a flat disk, especially at larger
radii. This model was able to reproduce the SED of T-Tauri stars without recurring to high
accretion rates or a combination of geometries; however, it failed to account for the near-infrared
3 Clm bump in the SED of HAeBe stars.
Dullemond, Dominik, and Natta (2001) proposed a modified Chiang-Goldreich model able
to fit the SED of HAeBe stars at all wavelengths. In a passively irradiated disk, dust very close to
the star is destroyed because of the very high temperatures in this region. As a result, the inner
disk close to the star is cleared of dust at the dust sublimation radius, truncating the disk at this
point. Gas might still be present in this "inner hole"; however, since the disk is passive, accretion
does not play a role and the holes remain optically thin to stellar radiation. The inner rim at the
dust sublimation radius directly receives radiation from the star and is very hot, as a result the
vertical scale height at this point increases casting a shadow over the disk behind it. The region
in the shadow becomes colder and collapses. At larger radii, depending on the opacity of the
disk, the disk can take two different geometries: if the optical depth is higher than a threshold
value, the disk' s outer regions flare and leave the shadowed region; if the optical depth is below
that value, the disk can never leave the shadow of the inner rim and remains flat and cold at
larger radii (Dullemond and Dominique 2004). This geometry, Figure 1-7, is able to reproduce
the SED of HAeBe stars and reconcile observations of these systems at different wavelengths.
The only drawback of the model is that it approximates the inner rim by a vertical wall and since
radiation from the inner rim is responsible for the 3 Clm feature in the SED, systems need to be
viewed at an inclination angle of ~45 degrees to have full view of the rim radiation (if viewed
close to face-on, the inner rim becomes a very thin annuli and the observed radiation from the
rim decreases considerably). This problem can be solved if curved inner walls in the rim are
introduced so that more of the rim radiation is visible for face-on disks as addressed by the
authors. This model will be addressed in more detail later in this research.
If the previously mentioned model can account for the observed SED of HAeBe stars
without recurring to an envelope, Vinkovic, Ivezic, and Miroshnichenko (2003) have shown that
by combining an optically thick flat disk and an optically thin halo, Figure 1-8, the same SED is
obtained. Since the dust in the flaring surface of the disk is optically thin, the temperature of this
dust depends only on the distance to the star and the grazing angle (the angle at which stellar
radiation is received) and spherical geometries with some assumed density profiles could
produce identical results. Therefore, it is impossible to distinguish between the two
configurations based only on the SED of the sources, leaving direct imaging the only way to
decipher the intrinsic geometry of HAeBe systems.
As we have seen, modeling the SED's of HAeBe stars does not provide unique physical
parameters to constrain the geometry. In addition, Brandner et al. 2002 showed that a slight
change in the viewing angle of a disk can lead to large differences in the observed SED of the
source (class I to class II) and large grains close to the star can mimic the spectra of small grains
further from the star. The detection of cold companions around some of these sources adds to
the complexity of HAeBe star environments. The uncertainties in the models and the lack of
consensus in the interpretation of the observations prove that, in order to make progress in the
study of PMS intermediate mass stars, direct observations of a large sample with enough spatial
resolution to see the inner regions around these sources and constrain grain properties of the CS
dust are needed.
This research focuses on the geometry and properties of the warm CS dust in HAeBe stars.
I observed 20 sources using the mid-infrared cameras T-ReCS and Michelle in the Gemini North
and South Telescopes to obtain high angular resolution and high sensitivity images in the 10 and
20 Clm regions. These observations allowed us to resolve or place severe limits on the sizes of the
mid-infrared emitting regions and characterize the structure of resolved sources. In Chapter 2, I
introduce mid-infrared astronomy and instrumentation and describe the observation technique
and data reduction procedures used in this study. In Chapter 3, I discuss our source sample and
data analysis. In Chapter 4, I present our results for AB Aurigae, the prototype HAeBe star. In
Chapter 5, I discuss the CS dust for Group I sources (in the classification of Meeus et al. 2001).
In Chapter 6, I describe the results for the Group II sources and discuss general results from the
survey; and in Chapter 7, I summarize the results and provide a coherent picture of the HAeBe
CS environments. I also discuss areas of future research.
11 12 13 1.4 15
Log v (Hz)
Figure 1-1. Stages of stellar and planetary system formation. Modified image from NASA
website. Image credits: Shu et al. 1987, Dullemond et al. 2001
Figure 1-2. H-R diagram showing the locus of different types of stars. Image credit: Pearson
Education, Addison Wesley.
Figure 1-3. Images of very young stars in the Taurus molecular cloud. The nebulosities seen
surrounding the stars are the result of small dust particles in the vicinity of the stars
reflecting stellar radiation. These images were taken by NASA Hubble Space
Telescope (HST) in the near-infrared using NICMOS.
~j20 -PAH -
2 diamond PAH.I
PAH Br | Sil icates
1P ~ I ~ II I I I
2 4 5 8 10 2
I PAH PAH .
2 4 6 8 10, 12
Figure 1-4. Infrared spectra of the HAeBe stars Elias 1 and HD97048. The emission features due
to PAH and silicate emission are labeld. Figure from Van Kerckhoven et al. 2002
9 01 1~ aI
1( 10 11
L~str + ~accetio (Lslar
Figue 15. ize-umiosiy coreltio dervedfro nea-inrard inerfromtricobsrvaion
0. 1 1.0 10.0 100.0 1000O.0
Figure 1-6. Typical spectral energy distribution (SED) of a HAeBe star. The figure shows the
stellar contribution and the 3 Clm bump.
no H2 yet
Figure 1-7. Schematic of the modified Chiang-Goldreich model developed by Dullemond (2001,
2004). The first panel shows the geometry of the dust in a flaring disk, the second
panel shows the geometry of the gas component. The gas component can extend to
the inner hole at the dust sublimation radius.
Figure 1-8. Schematic of the flat disk plus halo model described by Vinkovic et al. 2003 as an
alternative to the Dullemond model
MID-INFRARED INSTRUMENTATION, OB SERV NATIONS AND DATA REDUCTION
Infrared radiation covers the region from 1 to 350 Clm and is divided into three regimes:
near (1 to 5 Clm), mid (5 to 25 Clm), and far (25 to 350 Clm). Thermal emission is the principal
source of this radiation and any obj ect with a temperature above 0 degrees Kelvin emits in the
infrared. Infrared astronomy allows the exploration of many different sources: cool red giants
and star forming regions with temperatures in the 750 to 5000 Kelvin range in the near-infrared;
planets, comets, asteroids, and dust around stars with temperatures from 100 to 1000 Kelvin in
the mid-infrared; and very cold dust and molecular clouds with temperatures in the 10 to 100
Kelvin range in the far-infrared.
Many astronomical sources are only accessible to scientist in infrared wavelengths because
they are embedded in dense regions of gas and dust that obscure the obj ects at shorter
wavelengths, only becoming transparent at longer infrared wavelengths. In addition, obj ects that
are too cool and faint to be detected at visible wavelengths, like planets and dust in circumstellar
disks, can be detected in the infrared. Stellar radiation from stars of spectral type F or earlier will
heat optically thin dust to temperatures above 100 K within hundreds of AU from the star; this
warm dust will reradiate primarily in the mid-infrared part of the spectrum. At shorter
wavelengths, the disk becomes invisible because radiation from the star overwhelms and hides it;
however, at longer mid-infrared wavelengths radiation from the disk peaks, while radiation from
the star decreases, making it possible to detect the disk.
Ground-Based Mid-infrared Astronomy
The principal limitation in doing mid-infrared astronomy from the ground is the Earth's
atmosphere, which emits and absorbs infrared radiation. Thermal background emission from the
sky is primarily due to water vapor content, which is strongly related to ground temperature.
More water vapor leads to higher atmospheric emission. Terrestrial thermal emission, which
peaks in the mid-infrared, also contributes to the strong thermal background, limiting the
sensitivity of ground-based observations. The atmosphere also absorbs mid-infrared radiation;
Figure 2-1 shows the transmission curve from the atmosphere over these wavelengths. Regions
with fewer absorption lines, high atmospheric transmission, are called atmospheric windows and
most mid-infrared observations from the ground target these windows. In the 7-14 Clm window,
N band region, ozone is responsible for the strong absorption feature at 9.6 Clm, while water
vapor causes many of the absorption features in the 16-30 Clm window, Q band region.
Atmospheric transmission strongly depends on altitude; as altitude increases, the atmospheric
absorption bands become narrower due to a reduction in pressure, and emission from the sky is
slightly reduced due to a decrease in temperature. Therefore, good geographical sites for mid-
infrared observations are located at high altitudes.
Another maj or problem for mid-infrared astronomy is the unstable nature of the sky
background caused by microthermal turbulence and changes in water vapor and aerosol content
on short timescales. This effect is called seeing and causes blurring in the images affecting the
angular resolution of the observations. One way to measure the seeing is by measuring the full
width at half maximum (FWHM) of a point source. The image of a point source as seen through
a telescope is determined by diffraction. Light collected in the telescope lens interferes with itself
creating a ring diffraction pattern known as the Airy pattern. In the absence of atmosphere, the
angular resolution of the image is only limited by diffraction and given by the Rayleigh criterion
sin 8 = 1.22_ (2.1)
where B is the angular resolution in radians, Ai is the wavelength of the observation and D is the
diameter of the telescope lens, or telescope aperture. However, when light from an astronomical
source enters the atmosphere the parallel plane light waves are disturbed on small scales creating
an additional blurring effect on the images. This effect is greater for small wavelengths. If less
blurring is created by the seeing than is by diffraction on the telescope, the effect of seeing
becomes irrelevant and the resolution of the observations is only limited by telescope aperture.
However, mid-infrared observations from large ground-based telescopes are usually limited by
seeing conditions. Figure 2-2 shows changes in the FWHM of a point source during observations
at 11.7 Clm (first column) and 18.1 Clm (second column) taken from the Gemini 8 meter
telescopes. The first row shows observations taken under very stable seeing. The second row
shows observations taken under variable seeing.
The identical Gemini telescopes are two of the largest and more advanced optical
telescopes in the world and the first ones to be specially designed to work in the infrared. They
use a combination of two single mirrors in a Ritchey-Chritien configuration, Figure 2-3. This
optical design is used in all professional telescopes since it offers the advantage of eliminating
coma and spherical aberration, which can distort the image quality. The primary mirror is a
single 8.1 m mirror only 21 cm thick and is coated with a protective silver layer to reflect
infrared radiation more effectively while reducing the overall thermal emission of the telescope.
Since the mirror is too large and thin to keep its shape, it is tuned by using active optics at the
beginning of each night and every few minutes throughout the night. Tuning of the mirror is
done by splitting the light beam from a star and sending a representative part of the light to a
wavefront sensor that calculates how far out of shape the mirror is. Around the edge and under
the mirror, 180 computer-controlled actuators correct the deformation caused by gravity and
temperature changes by pushing and bending the mirror until it is perfectly shaped.
While adaptive optics is used to tune the primary mirror, adaptive optics is used to
minimize the effects of seeing in a process called fast guiding. A bright star near the science
target is used as a guide star. Light from the star is sent to a wavefront sensor, which calculates
the image distortion due to atmospheric seeing and the optimal position of the secondary mirror
to reduce it. The simplest form of this correction is using the tip-tilt mechanism of the 1-m
secondary mirror; the mirror position is adjusted around two axes so that the position of the
guiding star remains fixed in the field of view. Since the seeing can change rapidly during a
night, fast guiding continuously operates at a frequency greater than 10 Hz. The secondary
mirror of the telescope is also used for chopping during mid-infrared observations as will be
described in the next section, and is able to simultaneously do fast guiding and chopping.
Other features that make the Gemini Observatory optimized for infrared observations are
the excellent geographical sites and the system's thermal controls. Gemini North Telescope, also
known as Fred Gillet Telescope, sits on the dormant volcano Mauna Kea, Hawaii, at an elevation
of 13,822 ft. The Gemini South Telescope sits on Cerro Pachon, Chile, at an elevation of 8,895
ft. These are considered the best sites currently in use for optical and infrared observations due to
the stability of the seeing and atmospheric transparency. The system's thermal controls also help
improve the image quality by keeping the telescope, its enclosure and optics at the same
temperature as the environment, minimizing local self-induced seeing effects at the ground level.
Mid-infrared Instruments: Michelle and T-ReCS
After the telescope stabilize and correct the light beam using fast guiding, the beam is sent
through a central hole in the primary mirror to the instruments. Michelle and T-ReCS are the
mid-infrared imager/spectrometers in use at Gemini North and Gemini South respectively. Both
instruments are optimized for low thermal background and provide excellent image quality.
T-ReCS was designed and built for the Gemini South telescope at the University of
Florida. Figure 2-4 shows the path of the light beam inside the instrument and the different
* Entrance window: Environmental conditions, specially the percentage of humidity, and the
wavelength of the observation determine the entrance window used for an observation. The
entrance window fi1ter wheel provides three interchangeable entrance windows: (1) KBr
has a transmittance of 92% over the full mid-infrared wavelength range and emissivity of
only 0.04%, making it the ideal window for infrared observations; however, the window is
water soluble and not adequate for observing if humidity is high, (2) ZnSe has a
transmittance > 94 % and emissivity of 0.4%; this window is not water soluble but it has a
small wavelength coverage (2.5-13 Clm), and (3) KRS-5 has a transmittance of 70% over
the full mid-infrared range, an emissivity of 0.3% and is not water soluble.
* Filter wheel: T-ReCS has two filter wheels. Observations for this thesis were taken using
the Si-5, central wavelength 11.7 Clm, and Qa, central wavelength 18.3 Clm, filters. Figure
2-5 shows the transmittance pass-band for each of these filters.
* Detector: T-ReCs uses a Raytheon 320 x 240 Si:As block impurity band (BIB) detector
that covers the wavelength from 5 to 28 Clm. It is the largest mid-infrared detector currently
in use. The pixel size is 50 Clm. When used at the Gemini telescope, the pixel scale is 0.089
arcseconds for a total field of view of 28.4 x 21.1 arcseconds. The detector is made of 16
independent channels of dimensions (20 x 240 pixels). The signal of each channel is read
from the bottom in a consecutive order.
During operation, filter changes and telescope movement generate heat inside the
instrument. The temperature at which the array sensitivity is maximized is maintained using a
closed-loop temperature control system. T-ReCS is mostly a reflective instrument and all the
mirrors used are coated with gold to increase their reflectivity.
Michelle, the mid-infrared instrument currently in use at Gemini North, was built for
United Kingdom Infrared Telescope (UKIRT). The instrument has the same array as T-ReCS
and similar filters. The transmission band pass of the Si-5 and Qa filters used to collect the data
for this thesis are also shown in Figure 2-5. Michelle only has one entrance window made of
Potassium Bromide (KBr), so it is only used in low humidity.
Observing Technique: Chop-Nod
Despite all the care taken to optimize the instruments and telescopes for mid-infrared
observations, there is still higher background mid-infrared radiation compared to the radiation
we get from the science targets. To eliminate this excess emission, mostly due to the sky, and
locally to the telescope and entrance window to the instrument, a technique called Chop-Nod is
used. First the telescope is pointed at the target in the sky and a set of images is taken. The
images will contain mid-infrared radiation from the target superimposed on the large mid-
infrared radiation from the sky and radiation from the camera-telescope system. The frame time
of these observations is very small, ~10 ms; otherwise, the detector will saturate due to the high
background radiation. Then, the secondary mirror of the telescope moves slightly away from the
target source, and another set of images is taken. These images contain only mid-infrared
radiation from the sky in the vicinity of the source and from the telescope-camera system. The
telescope's secondary mirror will move from the target plus background position, or on-source
position, to the only background position, or off-source position, at a frequency of 3-10 Hz. This
procedure is called chopping. During a chop cycle mid-infrared radiation from the sky is
effectively removed from the images; however, because the position of the telescope' s secondary
mirror changes during the cycle, the path the light travels to reach the detector is slightly
different for the on-source and off-source positions and the radiation from the telescope-
instrument system is not effectively removed. The excess emission left after chopping is called
radiative offset. In order to remove radiative offset, the entire telescope is moved so that the
position of the source now appears where the only-sky position was previously, and the only-sky
position appears where the on-source position was. This movement of the entire telescope is
called nodding and occurs over a time frame larger than chopping, a few tens of seconds.
Position 1 is called nodA and position 2 is called nodB. In nod B position, the chop cycle is
repeated, and the radiative offset of the new difference images will be the negative of the one
obtained in nod position A. By combining the difference images from the two nod positions, the
radiative offset is effectively cancelled leaving only the mid-infrared emission from the target.
Figure 2-6 shows a standard chop-nod cycle. The figure was created using images obtained
with OSCIR at the Infrared Telescope Facility (IRTF). OSCIR is the University of Florida
infrared camera used to test the mid-infrared performance of the Gemini North and South
telescopes during commissioning of the telescopes.
Output Data Files Structure
The output data files for Michelle and T-ReCS are multi-extension Flexible Image
Transport System (FITS) files. The primary header, a text file containing all the parameters and
information used during the observation, is contained in extension 0 of the multi-extension fits
file. The remaining file structure is determined by which method is used to collect the data (the
chop-nod technique was used for the observations included in this thesis) and how the data is
saved in the two systems. Since the two instruments work in somewhat different ways, the
resulting file structure is slightly different. In Michelle output data files, each extension contains
the image data for a single nod position. All the images taken during a chop cycle at each nod
position are contained within an extension. As a result, the number of extensions in a Michelle
data file equals the number of nod positions used in the observation. In T-ReCS, each nod-cycle
is divided into multiple savesets of about 10 ms, which contain the co-added chopped frames
taken during that time. Each of the savesets is written into a different plane of the file extension
for that nod. There are usually 3 savesets per nod position. Therefore, T-ReCS raw data files can
be broken down into smaller sub-sets.
The nodding sequence for Michelle and T-ReCS observations is also different. During a
chop-nod cycle taken with Michelle, the telescope will move through positions ABBA. The
observations will start at nod position A, then go to nod position B, where two chop cycles will
be taken, returning afterwards to nod position A. During a T-ReCS chop-nod cycle, the
observations will start at nod position A, then move to nod position B, then move back to nod
position A, and then move back to nod position B. The sequence for a T-ReCS nod cycle is
Output Image Spurious Structure
The resulting image for each observation is a compilation of the many images produced
during the total number of chop-nod cycles taken during the observation. On occasions, the
resulting image will present distortions produced by non-perfect tuning of the primary mirror or
fast guiding of the secondary. The image can also present spurious structures related to the
detector. The data collected for this thesis was one of the first science observations taken with
the newly commissioned telescopes and instruments, and, as a result, the Observatory was still
correcting some of these problems. I only see image distortions due to telescope operations from
Gemini South observations taken in 2004. Figure 2-7 (A) shows some of these image
aberrations: the trifoil shape of the stellar image due to astigmatism caused by imperfect tuning
of the system, and an elongated structure across the image, probably also due to the telescope' s
Another problem with the output images is detector crosstalk. Fig 2-7 (A) also shows
spurious structures along rows and columns, which are characteristic of the detector response
when imaging very bright sources. The dark row centered on the source is due to a drop of the
signal along all channels and also shows ghost images of the source. The drop is only a few tens
of a percent of the peak emission of the stellar source and becomes obvious only in observations
of very bright sources. The image is presented in a logarithmic scale to emphasize the low level
structures. Since both Michelle and T-ReCS use the same detector, this image structure is present
at different degrees in all the observations. Since the effect is seen across the whole array and is
to some extent uniform, finding the median value along the rows for a range of columns with
only sky emission significantly reduces it.
The same method can be used to eliminate electronic noise in the images, which shows up
along columns and rows in the array. To clean the images, I modified an Interactive Data
Language (IDL) routine written by James M. De Buizer to create an image of the noise based on
the median value along columns and rows in the array. This image of the noise was then
subtracted from the images. Figure 2-7 (B) shows the resulting improved image after noise
In this section I describe the data reduction procedures used to prepare the output images
from the detector for scientific analysis. The following steps are standard for mid-infrared
2. Flat fielding
3. Sky subtraction
4. Standard star calibration and airmass correction
6. Color correction
Cross-correlation of the images, a useful technique to correct positional drift of the source
in the array due to tracking errors or imperfect fast guiding, was unnecessary for the data
included in this thesis. I looked at the data sets of all the images before stacking them into a final
image and measured the central position of the source in the array. Due to the excellent quality of
the Gemini Observatory, no drifts of the sources were observed in the images obtained for this
thesis and I was able to stack all the data subsets for each source to obtain a final image.
Flat fielding is another standard routine used in the processing of astronomical images and
corrects pixel-to-pixel variations in the detector array. The mid-infrared arrays used in Michelle
and T-ReCS are very flat, especially in the central region where the targets are centered. Since
taking flats for flat fielding involves more observing time and the non-uniformity in the array is
very small, the Gemini Observatory does not offer the option of doing flat fielding. Thus, no flat-
fielding correction was applied to the data.
Even after chop/nod there is sometimes residual background radiation in the images
making the mean radiation from the sky different from zero. Eliminating this residual emission
was the first step taken in reducing these data. The procedure used in all images is as follow:
* The image source was masked and the area in the perimeter of the source divided into four
regions (top, bottom, left, right)
* For each region the mean sky value (total count in the region divided by the number of
pixels in the region) and the standard deviation associated with the mean (sample standard
deviation, Upx, divided by the square root of the number of pixels) was calculated.
* The four sky regions were combined and a combined mean sky value calculated.
* The mean sky value was then subtracted from the image to effectively remove the residual
Standard Star Calibration and Airmass Correction
Photometric calibration is the process of converting Analog Digital Units (ADU/sec/px),
the flux units in the final detector image, into a physical flux density unit, in the case of this
thesis: Janskys, Jy. These observations were calibrated using mid-infrared standard stars. For
observations taken at Gemini North, the Gemini scientist overseeing the observations decided
which standard stars were observed on any given night. For some of the observations taken at
Gemini South, I was requested to provide specific standard stars to go with my observations. All
the standard stars were selected from two different lists of mid-infrared standards: (1) Mid-
infrared calibration standard stars compiled by the European Southern Observatory (ESO) with
accurate mid-infrared photometry from Thermal Infrared Multimode Instrument (TIMMI 2) and
(2) Cohen et al. (1999) mid-infrared standards. In the cases when Cohen (1999) standard stars
were used, I calculated the flux densities for these star from the spectral models provided by
Cohen. Only two observations of standard stars at each filter were taken along with each
Once I know the expected flux density of the standard at the filter central wavelength, I
need to find calibration values. Since the mean sky is already zero, this is a straightforward
procedure. I used a circular aperture to measure the ADU/sec/px in the image of the standard star
and divided the expected flux by that value; the ratio of the two is the calibration value for that
observation. The calibration value is an indicator of the photometric quality of the night. The
error associated with the flux measurement in the images, Ox where npx, is the number of
pixels included in the aperture, is much less than the 10% standard error associated with
photometric calibration in the mid-infrared; therefore, I assume the typical 10% error for the
Under ideal conditions, the calibration value will be the same for all stars observed at the
same airma~ss. Airmass is a measure of the thickness of the atmosphere that light from an
astronomical source has to cross to get to us relative to the zenith and is defined by
Airmass=sec(z), where z is the angle of the observation from zenith. To reduce the negative
effect of the atmosphere on the observations, it is desirable to take observations at an airmass
close to 1. All the sources in this survey were observed at airmass < 1.5. In reality, the airmass
correction needed in the mid-infrared is very small. To include airmass correction, standard stars
are observed at different airmasses throughout the night with each fi1ter so that the airmass of the
science target can be matched to the airmass of a standard star. Since this procedure adds a
significant amount of time to each science target observation, airmass correction was not applied
to the images. In all these observations, the calibration values obtained from the two standard
star observations at each filter were almost the same.
Once the calibration value for one observation is known, I need to multiply the science
target image by that value so that the flux density is in Jy/px. A circular aperture is placed over
the source to measure the total emission from the source. The size of the aperture was decided on
upon inspection of each image to include all flux from the source and as little sky as possible.
The error associated with this measurement is again dominated by the 10% error in the
Color correction is important for flux density measurements, because the SED of stellar
sources used as standards and the SED of the science targets have different shapes and therefore
different colors over the fi1ter bandpass. Color correction depends not only on the shape of the
SED of the two objects, but also on the central wavelength and bandpass of the fi1ters. The effect
is larger for wide band filters. In this thesis, only narrow band filters were used; however, there is
still a small error associated with color-correction.
Most of the standards used in this thesis are spectral type K or M with effective
temperatures in the range of 3000 to 5000 K. These sources were selected to avoid using the very
bright stars that saturate the detector; however the AO star Sirius with a temperature of 9,500 K
was also used in some observations. The SEDs of these sources peak in the optical and have a
steep decline dominated by the Rayleigh-Jeans law, f(il) ac il2, in the mid-infrared. On the
other hand, the SEDs of the target sources are very flat at this wavelengths because they are
dominated by grain dust emission with temperatures of ~ 100-300 K that peak in the mid-
The color correction method assumes that
where P stands for photometric standard, S for science target, and Nr is the number of photons
received at the detector. The number of photons is given by
N = Q jdV (2.3)
where vl and v2 are the frequency limits of the observation, h is the Planck' s constant and F, is
the flux density of the source through the filter bandpass. The flux density depends on the solid
angle of the observation, Dn, the optical depth, rv, and the Planck function as a function of
temperature, B, (T),
F, = O2(1 e )B, (T) (2.4)
and Qsys is the system quantum efficiency which depends on the transmission of the atmosphere,
Tarran, and the filter, Tyurer,, and on the quantum efficiency, QE, of the detector at the desired
Gas Tru, (A Taran A)QE(2)(2.5)
Combining the previous equations, I can write the number of photons in terms of a
monochromatic flux density Fvo
F, 2z SPSe ~B, (T)
N = o d (.6
(1- e-rvo)Bv(T) h ss V ~h 26
and relate the number of photons received from the program source to the counts measured in
ADU by using the first equation in this section. Approximating the standard star by a blackbody,
the opacity term for that side of the equation disappears. Then I can write the monochromatic
flux density for the science source as a function of the ratio of the flux density to the counts
measured for the standard star, which is the calibration value, multiplied by the counts in ADU
measured for the source, and a color correction term. The calibration value multiplied by the
source counts is the calculated flux density for the source and the color correction from the
previous derivation is
Q B,, (~T) dV
color correction term = vB T)(2.7)
S (1- e ')B,(TP) d
Color correction was done using a MATHCAD program written by Dr. R. K. Pifia. The
input parameters are the calculated flux density and filter properties. The code iterates to solve
for temperature, opacity and color correction. Since the filters used in this thesis are narrow band
filters, the color correction for all sources using the previously outlined procedure is very small,
less than 1 % for the Si-5 filter and less than 4% for the Qa filter. Therefore, the color corrected
flux measurements are within the photometric errors.
5 10 15 20 25
Figure 2-1. ATRAN model of atmospheric transmission for Mauna Kea. The absorption feature
at 9.8 Clm is due to ozone.
0.6 0.8 1.0 1,2 1.4 1,6 1.1
n > I .
7.0 7.5 8.0L 8.5 9.4
as I I ,
Figure 2-2. Full width at half maximum measurements as a function of Universal Time. The first
row shows observations taken under stable seeing. The second row shows
observations taken under variable seeing. The 18.3 Clm observations are less affected
by seeing because of the larger wavelength. Both sets of observations were taken as
part of this thesis data at Gemini South telescope using T-ReCS.
Figure 2-3. Optical design for a Ritchney-Chretien telescope.
Figure 2-4. T-ReCs optical design
0.8 T-ReCS 0.B
0 .4 0.4-
0.2 kee0.2 .TRC
10.5 11 11.5 12 12.5 13 16.5 17 17.5 18 18.5 19 19.5 20
wavlength (pn0) walength (pa0
Figure 2-5. Filter transmission for T-ReCS (blue) and Michelle (magenta)
Nod Posi-tion A
Nod Posi-tion B
un-hource ...' ut-bource
Chopp ed Difference "Chopped Difference"
Net Source Signal
Figure 2-6. Chop-Nod strategy. The top four images are the raw co-added images at each chop
position. The target source is not seen on the on-source images because emission
from the source is small compared to sky emission, which dominates the mid-infrared
radiation observed. The middle row shows the resulting images after subtraction of
the off-source images, removing the contribution from the sky. These two images,
called dif frames, are the negative of each other and show the large radiative offset
due to the telescope-camera mid-infrared emission. The bottom image is the result of
combining the two dif frames and mid-infrared radiation from the astronomical
source dominates the final image. This image was taken with OSCIR at the IRTF and
is part of the OSCIR operation manual.
Figure 2-7. Spurious structure in the images. The image on the left shows image structure due to
diffraction (airy ring), imperfect tuning (trifoil structure and diagonal line across the
source), detector response to bright sources (ghost image across array), and electronic
noise (horizontal banding noise). The image on the right shows the improved image
after implementing the noise mask routine. Non-uniform structures are still present.
SOURCE SAMPLE AND DATA ANALYSIS
Herbig Ae/Be Sample
We selected all targets for this survey (except one, HD97048) from the list of optically
visible Herbig stars by Malfait et al. (1998). The sources were selected by matching the positions
of stars in the Smithsonian Astrophysical Observatory (SAO) Star Catalogue and stars in the
IRAS point source catalogue. An advantage of this selection is that this sample should not
include any of the heavily obscured very young systems and I can concentrate on the sources that
have evolved beyond the embedded stage. A drawback of the selection criteria is that the sample
also excludes stars with disks orientations far from face-on where the star is obscured by dust in
the disks, and the appearance of face-on disks is difficult to distinguish from spherical halos. All
20 obj ects were chosen to be within 500 pc, with 17/20 obj ects being within 300 pc, to ensure the
ability to spatially resolve the circumstellar structures in these systems. Many of the sources in
the target list have been shown to have extended CS dust from studies at other wavelengths.
Calculating Stellar Properties
Determination of stellar parameters and evolutionary states is limited mostly by the
accuracy of stellar models. The morphology of the stellar tracks from different stellar tracks
codes (Baraffe et al. 1998, BCAH code; Charbonnel et al. 1999, Geneva code; Siess et al. 2000,
Grenoble code) largely depends on the different physics, convection, surface boundary
conditions, that goes into the codes. However, for stellar masses above 1 solar mass and ages of
more than 1 million year, the regime of interest for us, all the stellar tracks are very close to each
Spectral types, effective temperatures and luminosities of all the target sources were taken
from the literature (Acke et al. 2005, van den Ancker et al. 1998, van Boekel et al. 2005, Pietu et
al. 2003, Jura et al. 2001, Hamaguchi et al. 2005, and Moth et al. 1997). I then used those stellar
parameters and assumed solar metallicity (z=0.02) to derived stellar masses and ages for all the
sources in the survey by using the stellar PMS evolutionary tracks of Siess et al. (2000). To
check the accuracy of the calculated stellar parameters, I compared the derived parameters for
the very well studied Herbig Ae star AB Aurigae to other published results. The parameters I
derived were within the range of published values found in the literature for this source. An
overview of the stars in the sample, some of their properties, and calculated stellar parameters
are given in Table 3-1.
An H-R diagram for all the sources, Figure 3-1, shows the maj ority of the sources along
evolutionary tracks for PMS stars with masses between 1.7 and 7 solar masses. Most of the stars
in this study are approaching the zero age main sequence line, reinforcing the idea that they have
probably dissipated most of the original envelope material. However, the classification of two of
the sources as young stars and therefore candidate Herbig Be stars is questionable. The star
HD41511 is thought to be a binary system with an M type companion star and its position in the
H-R diagram well above the main sequence has been interpreted as evidence of its post-main-
sequence state (Welty and Wade 1995). The star HD50138 is also a suspected binary system
(Cidale et al. 2001), and it has been classified in multiple studies as a classical B[e] star.
Flux densities measured from the images contain emission from the stellar photosphere
and thermal emission from dust grains around the star. Since the main objective of the study is
to investigate the properties of the mid-IR emission from the dust, I need to isolate this emission
by removing the photospheric component of the emission.
Two different methods were used to estimate the stellar contribution to the source fluxes
derived from the observations. First, I used the spectral type of the star and its effective
temperature from table 3-1 to select the appropriate stellar atmosphere model using the
PHOENIX/NextGen grid of models developed by Peter H. Hauschildt for M dwarf stars of solar
metallicity. This model includes over 500 million atomic and molecular lines and treats stellar
structure as a sphere instead of using the parallel plane approximation. As a simpler alternative, I
used the effective temperature of the sources and approximated the stellar photosphere by a
blackbody function. Since I was only interested in the mid-infrared contribution, the blackbody
approximation also gives an accurate estimate of the stellar flux. I then used photometric
measurements in the optical UBV bands to scale the models. The average ratios of the Uo/Um,
Bo/Bm, Von/m, where o stands for observed and m for modeled, were used to calculate the scale
value and normalize the model atmosphere.
From the normalized atmospheric model, I found the expected atmospheric fluxes at the
wavelengths of the observations. These fluxes were then subtracted from the total fluxes to
obtain the dust emission contribution. Figure 3-2 shows the stellar atmospheric model and the
blackbody approximation, both normalized to the UBV fluxes for one of the sources. The
derived mid-infrared fluxes are identical using either of the two approximations. Since I did not
have a theoretical computation model atmosphere matching the effective temperature of all the
sources from NextGen, I only used photospheric fluxes derived from the blackbody
approximation when the NextGen model was not available.
For all the sources studied in this survey, stellar fluxes were negligible when compared to
the total flux of the system. Mid-infrared emission in these obj ects is dominated by dust
In addition to correct the fluxes for photospheric emission, I also corrected the 11.7 Clm
fluxes for silicate and PAH line emission. I used the available line fluxes from Acke and van den
Ancker (2005) to estimate the contribution of the 1 1.2 Clm PAH emission line and the 1 1.3 Clm
silicate emission band to the fluxes calculated from the observations since these emission lines
are within the Si-5 filter bandpass. For all the sources with line flux information, the line flux
was subtracted to leave only the contribution from dust thermal emission at this wavelength.
Classification of sources
For consistency with the literature I have classified all the sources into the Group I/Group
II classes as defined by Meeus et al. (2001). I used near-infrared photometry in the JHKLM
fi1ters (Malfait et al. 1998, Hillenbrand et al. 1992) to calculate near-infrared luminosities, and
flux densities at 11.7 and 18.3 Clm from the observations to calculate mid-infrared luminosities
and mid-infrared colors -. This differs from previous classifications in that I am using
the mid-infrared fluxes derived from the observations presented in this work and not IRAS
fluxes. The resulting graph is shown in Figure 3-3. Only two sources have different
classifications than in the literature. HD34282 has been previously classified as a Group I source;
however, I classify this obj ect as a Group II source based on the photometry presented here. The
star HD 150193 also falls under a different classification. It has been previously classified as a
Group II source. Based on the fluxes derived from the observations presented here, HD150193
would be a Group I obj ect. However, the Si-5 observations of this source were taken under
unstable conditions, with seeing variations as large as 0.4 arcseconds, and I cannot evaluate the
photometric errors during that night since only one observation of a standard star was taken.
Given the uncertainties associated with the Si-5 photometry of this obj ect and the lack of
nebulosity around the source that could compromise the IRAS measurements, I have kept the
Group II classification for HD 150193.
The accurate measurement of angular diameters of partially resolved sources is a long-
standing problem. Nevertheless, very few papers have been published on this subj ect. In this
section I describe the systematic approach I followed to measure sizes of the target sources.
Point Spread Function
The image of a point source on the detector is not a point; because of diffraction the
observed image is a convolution of the intrinsic source and the instrumental profile or point
spread function (PSF). The PSF can be determined by observing a main sequence star close to
the program star on the sky before and after each observation of the program star. It is important
for the observations of the PSF and science star to be close in the sky and in time so that
observing conditions for the two are as similar as possible. Since these main sequence stars are
point sources without extended structure, the resulting image profile is a reflection of the
instrumental profile and seeing conditions relevant to the program star.
All PSF stars for the survey were selected from the Positions and Proper Motions (PPM)
Star Catalogue. I constrained the list to include only stars of spectral types K or M, brighter than
V = 6, and located within 5 degrees of the program star in the sky. One of the problems I
encountered early on with the observations of PSF stars was that although they were bright at
visible wavelengths, they were faint at mid-infrared wavelengths. Since observations of PSF
stars are taken with short integration times, the resulting PSF image had low signal to noise ratio
for the desired analysis. In an effort to correct this, all PSF stars for subsequent observing runs
were correlated with IRAS point sources. Selected sources from the PPM catalogue with no flux
measurement in the IRAS catalogue or with low flux density at 12 and 25 Clm were discarded.
Very bright mid-infrared sources were also discarded to avoid image distortions due to detector
crosstalk. This step significantly improved the PSF star observations.
Another blurring effect on the images that I have not mentioned previously is due to pixels
on the array. The image registered on the detector is broken down into pixels and each pixel act
as a boxcar function so that the signal from each pixel is a constant. The resulting effect can be
seen if one zooms in the image; instead of having a smooth source profile, one finds sharp edges.
For slowly changing regions of the source profile, interpolation can be used to break the original
pixels down into smaller pixels. New data points are created by smoothly blending the signal
from pixel to pixel. However, pixelation is more difficult to correct for sharp regions of the
profile, with only one or two pixels covering the whole region. In this case, an alternative
method is to model the source profile using a well behaved mathematical function.
Two widely used functions to model stellar profiles in astronomy are the Gaussian function
and the Moffat function (Moffat 1969). The Moffat function provides a better fit to source
profiles especially on the wings of the profile, where a Gaussian function decays faster. The
Moffat function is given by
I(r) = 1 + (.1
and in the limiting case where P 4 00, it is identical to a Gaussian function. Figure 3-4 shows
the normalized profile of a Gaussian and Moffat functions with different values, the smaller the
p value, the larger the wings of the profile.
The presence of wings in all the sources makes the Moffat function an excellent choice to
model PSF and science targets, and recover the stellar profiles without the effect of pixelation. I
used the IDL code mpfit2dpeak created by Dr. Craig B. Markwardt and publicly available at
(http:.//cow.physics.wisc. edu/~craigm/idl/fitting.html) to model the profiles.
In order to extract meaningful sizes for the system, it is necessary to deconvolve the
observed image and the PSF, especially if the profiles of the science target and PSF star are
similar. The problem with deconvolution is that it is a complicated technique and the results
strongly depend on the assumptions made and the method employed. As a simpler alternative, I
have chosen to deconvolve the observed sizes at the FWHM by using quadrature subtraction of
the observed source diameter, 0, and PSF diameter, GQ. The deconvolved diameter, ed, and the
error in the size, 00d, are then given by:
Od 2_ b (3.2)
At the FWHM level, all the sources diameters are very close to the PSF diameters. Since I
integrated longer on the sources than on PSF stars and integration time can influence the image
profile, especially due to changes in the seeing during the observation, I divided all the images
for HAeBe stars and PSF stars into multiple subsets of equal integration time. This offers the
advantage of excluding integration time as a factor for wider FWHM in HAeBe stars, while
providing information about the seeing conditions during the observation and a good statistical
sample to strengthen the results. This technique could only be used to measure diameters at this
level, where the surface brightness of the individual data frames was well above the 30 level of
We considered a source to be extended if the deconvolved diameter minus the error was
larger than zero. For all unresolved sources, I calculated limiting sizes at the FWHM brightness
contour level by increasing the FWHM of the HAeBe stars until the deconvolved size minus the
error was larger than zero.
A last step taken to evaluate extension in the systems was to subtract the PSF from the
sources. First, I normalized all images so that the peak fluxes in all HAeBe stars and PSF stars
equal 100. This was done by dividing all the HAeBe and PSF star images by the peak fluxes in
their images and multiplying by 100. I then cropped the images so that the peak position for PSF
and HAeBe stars were on the same position in the images, and subtracted the PSF image from
the HAeBe star image. Residual emission after PSF subtraction is another evidence of extended
emission. To measure diameters of this more extended fainter emission, I chose surface
brightness contour levels equivalent to 0.5 % and 1 % of the peak emission at 11.7 Clm, and 5%
and 3 % of the peak emission at 18.5 Clm. These specific levels were chosen so that they were
more than three times the background noise for all extended sources and their respective PSF
stars. Since I used stacked images in this procedure, all observations without stable seeing, as
reflected in the FWHM measurements, were discarded from this analysis. In order to assess the
validity of the residuals, I directly measured the diameters at the specific brightness levels from
the images and used half the spread in sizes observed in the FWHM as an upper limit to the
Properties of the Circumstellar Dust
In this section I explain the different analysis I used to derive properties of the
circumstellar dust for all the sources in the survey.
The temp We directly measure the diameters at this level from the images. erature of a dust
grain in radiative equilibrium depends on the optical properties of the grain and its distance from
the radiating source.
The flux from an isotropic source at a point P at a distance r from a star of radius R and
effective temperature Tefis
F = B (T,)cosMD 3.4
where daY = 2xisin MO is the solid angle of the observations. From geometry, I can then write:
F,= ,(Tjsn20 m,(T, (3.5)
Then the power from the central star absorbed by a uniformly heated dust grain at that
location will be
where a is the dust grain radius and Q,, is the emissivity of the material, which also depends on
grain size; grains absorb and emit radiation more efficiently for wavelengths smaller than the
The power radiated by the grain is then given in terms of the dust grain temperature, T,, by
F (rad) = 4Wu Q mS (T)d v 37
and, given the condition of radiative equilibrium, I can then write
SB, (T,,)Q,du =4 B, (T, )Q, dv (3.8)
with the assumption:
where Alo is a critical wavelength in the sense that absorption and emission of radiation at
wavelengths shorter than Alo gives Q,,~ 1.
Since the absorbed and the emitted radiation have different wavelengths, I can let n=p for
absorbed radiation and n= q for emitted radiation. Then I can use different exponents, p and q, to
approximate the behavior of different grain materials. Grains much larger than incoming and
emitted radiation can be approximated as blackbody grains with p=q=0; grains much larger than
the incoming radiation, but smaller than the peak of radiated emission will have p=0 and q=1 for
amorphous materials and q=2 for crystalline materials and metals; grains smaller than the
incoming and emitted radiation like ISM grains will emit and absorb inefficiently; thus, I can
assume that p=q=1.5 (Helou 1989, Backman and Paresce 1993).
Therefore, given some assumptions about the composition of the dust grains, I can solve
Equation 3.8 analytically to obtain the temperature of the dust grains. Substituting for the Planck
where c is the speed of light, h is the Planck constant, and k is the Boltzmann' s constant into
Equation 3.8 and introducing the Stefan-Boltzman constant, cr = I have:
c (Rif 4-Lu21 v" exp- -1 dv] (3.11)
rf V4 ,c k T
By letting x = we can solve the integral analytically and obtain:
U~q,4 8(kT,)q+4h (4 j)(q+3)Z(q+4) (3.12)
Defining the stellar luminosity by:
L, = 418R UT 4~ (3.13)
we can then solve for the temperature of grains of different materials as a function of distance to
Tq+ ,,Ch+ (3.14)
32nQ,k q4'4(q + 3)Z(q + 4)r
For blackbody grains, where q=0, this is the equivalent expression:
T,= 278L, r ,l (3.11)
For medium size grains with q=1 I get:
Tg = 468L*5r A,~ (3.12)
and for very small grains, where q=p=1.5, a different expression has to be used, which
Ts = 636L."~r "Tq, (3.13)
Warm Dust Mass
We also calculated the warm dust masses in these systems. Using the dust temperatures
derived from the flux ratios and the dust thermal emission at 18.3 Clm, I can provide a rough
estimate of the minimum warm dust mass contributing to the observed emission by using,
n"B, (T, )X,
where Fris the flux density from dust thermal emission at 18.3 Clm, Tg is the derived average
dust grain temperatures for the systems, D is the distance to the sources in parsecs as listed in
Table 3-1, and X, is dust absorption opacity estimated to be 1000 cm2/g (Ossenkopf, Hennings &
Dust Size Distribution
I used an approach similar to the one explained in the Dust Temperature section to derived
typical dust grain sizes. Instead of plotting the temperature of the grains as a function of radius, I
plotted the expected flux ratios for particles of different sizes and values of ilo as a function of the
distance to the star. I then used the flux ratio from the observations and assigned this value to the
mid-point of the resolved extended emission region or the mid-point of the limiting size in the
case of unresolved sources. Just by inspection, I can find which modeled flux ratio better
represents the observed flux value (large blackbody grains, very small grains, or mid-size grains
of some specific ilo). In the case of mid-size grains, the value of Alo can be used to estimate grain
sizes: a ~ ilo for moderately absorbing dielectrics and a ~ io/27r for strongly absorbing materials
State of Evolution of Dust Grains
Studying the effects of the ratio of the radiative to gravitational forces, P, on the
maximum sizes derived above, I can assess the evolution of the dust in these systems.
where c is the speed of light, G is the gravitational constant, Leis the stellar luminosity from
Table 3-1, Ms is the stellar mass also from Table 3-1, a is the size of dust grains derived above,
and p is the dust density. For P > 1, the radiative force will predominate. Studies of the solar
system zodiacal cloud have shown that for realistic grain materials, particles are in bound orbits
forp< 0.5. All particle sizes in a system giving P> 0.5 will be in unbound orbits and expelled
from the system in a timescale comparable to the orbital period.
Stellar optical radiation penetrates the CS dust to a radius corresponding roughly to an
optical depth zv = 1. If I assume a flat disk, with inner radius, Rln equal to the dust sublimation
radius, outer radius, Ro,,,,equal to the larger disk size known for a system, thickness AS, a cold
dust mass, M,,,,,, from sub-mm and mm studies when available, and uniform volume dust po
density given by the relationship:
Po =z(R,,, RjAS(3.16)
we can calculate the optical depth along the plane of the disk by using:
r, = rl, pdr (3.17)
with the absorption coefficient of the form
r, = r,(3.18)
where rno = 0. 1 cm12 g-1 (Hildebrand 1983). From these calculations I can estimate the radius of
the disk for which z, = 1, where the disk becomes optically thick. This value is an upper limit in
the sense that the inner regions of a flaring disk are likely to be thinner than 10 AU, and any
other axisymmetric dust density distribution will place more dust mass closer to the star making
the disk opacity higher at small radii. Knowing where z,=1 for a system tells us how far stellar
radiation penetrates into the plane of the disk. Dust located at larger radii should be cold and;
therefore, undetected at mid-infrared wavelengths. If I detect warm dust beyond this radius in a
system, it implies that the dust must be heated by direct radiation from the star and consequently
should be located above or below the disk mid-plane in an optically thin region either a halo or
the surface layer of a flaring disk.
To further constrain the location of the mid-infrared emitting dust in the system, I
considered the vertical (i.e. perpendicular to the plane of the disk) optical depth in a system.
Given the low optical extinction toward the central star, I believe most of these systems to be
close to face-on; therefore, an estimate of the vertical optical depth of the disk at mid-infrared
wavelengths should indicate if I are only detecting emission from the dust located above the disk
(if the disk is optically thick to mid-infrared radiation) or from both sides of the disk (if the disk
is optically thin to mid-infrared radiation). The disk in the mid-plane evidently would be too cool
for us to detect even if it were optically thin to mid-infrared radiation since optical radiation
cannot penetrate there.
For a disk surface density with the radial distribution
r = 1 (3.19)
where Zois the disk surface density in g cm-2 and
r, = inCE (3.20)
we can calculate how far from the central star the disk becomes vertically thin to mid-infrared
Passive Flared Disk Model
The Chiang and Goldreich model is based on the idea that the surface of a CS disk is
irradiated by the central star under a shallow angle. Previously in this chapter, I derived the
energy balance equations for an optically thin disk. A grain of dust in a thin disk sees radiation
from all the stellar surface facing it. This is not the case for very massive, optically thick disks.
The previously derived equations have to be modified then to include the angle oc of incidence of
stellar radiation upon the disk, (grazing angle). The irradiated flux can then be written by
F~ = a (3.21)
and the emitted flux disk temperature takes the form
T ~ T,f (3.22)
For a flat disk, the scale height, H, is independent of radius, and the grazing angle, a, and
can be approximated by as for r >> R (Ruden and Pollack 1991). Chiang and Goldreich
(1997) presented a more general form of this equation accounting for the height H of the disk
surface above the mid-plane:
0.4R d HIH (3 H3
a +r-- 4-(.3
r rr r
The irradiated dust at the surface of the disk absorbs stellar flux and re-radiates in the
infrared part of the spectrum. About half of the infrared flux is radiated away from the disk
plane and escapes the system; the other half is sent downwards and heats the dust grains in the
disk' s mid-plane. The higher densities expected in these more massive optically thick disks
translate into thermal coupling of the dust and gas in the disk (the gas and dust in the disk have
equal temperatures). The heating of the mid-plane by the re-radiated infrared emission from the
optically thin outer layer heats the dust and gas in the mid-plane, resulting in higher gas
temperatures and increasing the gas pressure. The increase in gas pressure counteracts the effect
of the gravitational forces pulling disk matter towards the mid-plane and affects the vertical scale
height, H, of the disk and, consequently, changes the geometry of the disk.
The Chiang and Goldreich model assumes thermal and hydrostatic equilibrium, thermal
coupling of the dust and gas, a gas to dust ratio of 100 and a form for the opacity of dust grains
and surface density distribution. Given those assumptions and the luminosity of star, the model
calculates the expected surface temperature of the disk. Using the previous equations and writing
the effective temperature in terms of the stellar luminosity, the surface temperature is given by
T4 = ( (3.24)
From the surface temperature, the model derives the flux irradiated inward towards the
disk mid-plane and calculates the temperature and pressure of the gas in the disk' s mid-plane.
Using the balance between gravity and gas pressure, the model is solved for hydrostatic
equilibrium by iterating from a flat disk geometry and adjusting the scale height of the disk until
balance is reached.
The surface scale height H can be written in terms of the pressure scale height, H = Xh,
where X equals a constant and the pressure scale height is defined by
h = (3.25)
where pg, is the mean molecular weight of the gas.
From Equation 3.24 I can write
and obtain the relation between pressure scale height and radius, h ac r9/7 The resulting
geometry is a scale height that increases with radius, a flaring geometry, Figure 3-6.
The modified Chiang and Goldreich model (Natta et al. 2001, Dullemond, Dominik and
Natta 2001) starts by clearing the dust inside the dust sublimation radius. Dust can still be
present within this region; however, because I are assuming a passive disk, this region is
optically thin to stellar radiation. The result of truncating the disk is that dust located at this
radius receives stellar direction upfront, at an angle of 90 degrees, and heats up more than in the
Chiang and Goldreich model with a low grazing angle. Because of hydrostatic equilibrium in this
region, the vertical scale height of the disk at this inner radius becomes large and creates a wall.
The larger vertical scale height implies larger emission surface, which explains the excess near-
infrared bump at 3 Clm in the spectral energy distribution of HAeBe stars. The scale height of the
inner rim is computed as the surface scale height of the flaring disk, but using a grazing angle of
The puffed-up inner rim affects the geometry of the rest of the disk. It intercepts more
stellar radiation and the disk region directly behind the rim lies in its shadow and remains cold.
This shadowed region is heated only by the fraction of stellar light scattered towards the disk by
the edge of the rim and by thermal diffusion from the inner rim, resulting in a very cold region
with little contribution to the total emitted radiation from the disk. The geometry in the outer
regions of the disk can take two solutions. If the disk opacity is high, the surfaces of the outer
parts are still irradiated as in the Chiang and Goldreich model. Re-emitted stellar flux from the
inner rim and surface layer towards the disk' s mid-plane maintains the temperature; therefore,
the pressure of the gas high enough to increase the scale height of the disk, resulting a flaring
region. In the second solution, the geometry of the system in the outer regions remains at low
scale height. This is the result of low dust opacities; the absorbed flux is not sufficient to push
the surface layer up and the disk stays in the shadow of the inner rim. Figure 3-7 shows a
schematic picture of these two geometries and the expected SED for each case.
Table 3-1. Physical properties of program stars.
HD # Sp. Type d [pc] Log Teff [K] log L [Lsun] Mass [Msun] age [Myr]
31293 AOV 144 3.98 1.71 2.57 3.81
34282 AOVe 400 3.94 1.38 2.19 6.08
36112 A3-51VIV 204 3.91 1.47 2.00 5.67
36917 AOV 460 3.93 2.20 3.72 1.22
37806 B9-A2V 470 3.95 2.13 3.36 1.69
38087 B5V 470 4.20 2.56 4.98 2.03
41511 A2esh 328 3.97 3.10 7.00 3.94
50138 B5-B9Ve 290 4.19 2.85 5.02 0.79
58647 B91V 280 4.03 2.48 4.06 0.80
97048 B9-AOV 180 4.00 1.61 2.53 4.96
100453 A9V 112 3.87 1.09 1.84 14.69
135344 AO-F4V 140 3.82 1.01 1.74 9.70
139614 A7Ve 140 3.89 1.03 1.85 10.45
144668 A5-71ll/IV 210 3.90 1.98 3.26 1.65
150193 Al-21VIV 150 3.95 1.38 2.18 7.08
158643 B9-AOV 131 4.00 2.39 3.88 1.22
163296 AO-A31VIV 122 3.94 1.52 2.30 4.56
169142 B9-A5V 145 3.91 1.13 1.94 10.59
179218 B9Ve 240 4.00 2.50 4.33 0.73
244604 AO-A3V 336 3.95 1.35 2.18 7.53
Note: columns (2-5) Acke et al 2005, van den Ancker 1998, van Boekel 2004, Pietu et al. 2003,
Jura et al. 2001, Hamaguchi et al. 2005, Moth et al. 1997. Columns (6, 7) calculated from stellar
PMS tracks (Seiss et al. 2000)
4.4 ~~4.2 4 .83.
Fiue -. dara oralth oucs nlue i hi uve.PM vouioay rck r
10 4 I I I.,. ...I ....., I 1 .....,,, ,, I ..~, ,
102 10" 104 105 10e 105
Figure 3-2. Spectrum from stellar photosphere using the NextGen model is shown with the solid
line. The dashed line shows the blackbody approximation for a stellar source with the
same temperature. The two curves are very similar in the infrared regime of the
-0 0 I -Ii---- -
| | Group I
-2 -1 0 1 2
Figure 3-3. Plot of the ratio of near- to mid-infrared luminosity versus mid-infrared colors. The
open triangles represent Group II sources, while the black filled circles represent
Group I sources in the classification of Meeus et al. (2001). The open circle is
HD 150193 previously classified as Group II source and the filled triangle is
HD34282 previously classified as Group I source. The red dotted line divides the two
Figure 3-4. Normalized intensity profile of a Gaussian function and Moffat functions with
different parameters P. The difference in the profiles is in the wings. The lower the P
value, the larger are the wings of the profile.
ujillo et al. 2001
- - -P= 2.50
---- --- =4.76
hot surface layer
Figure 3-5. Representation of the Chiang and Goldreich flaring disk. The dust disk extends all
the way to the central star. The optically thin surface layer of the disk receives stellar
radiation at a low grazing angle, while the mid-plane remains colder.
Figure 3-6. The two solutions for the modified Chiang and Goldreich disk model. The inner
region of the disk is cleared because the dust is destroyed at temperatures above 1500
K. The dust at the edge of the disk is then heated directly by stellar radiation and its
scale height increases because of gas pressure. The area directly behind the puffed-up
inner rim stays in the shadow of the rim and is very cold. At larger radii, if the optical
depth of the disk is high, the disk will assume the flaring disk geometry of the Chiang
and Goldreich model; however, if the optical depth of the disk is low, the disk will
remain flatter and in the shadow of the inner rim. These two geometries have been
associated with the Meeus et al. (2001) classification of Herbig Ae/Be stars in Group
I (flaring disks) and Group II (self-shadowed disks).
Located only 144 pc away (van den Ancker et al. 1998), the 2 Myr old AO star AB Aurigae
is the brightest (V = 7.06) of the original sample of Herbig stars (Herbig 1960), which are
intermediate mass (2 to 8 Me) pre-main sequence stars. Consequently, it is not only the best-
studied Herbig obj ect, but an important touchstone for the understanding of the class. The
spectral energy distribution (SED) of this source shows emission in excess of the photosphere
throughout the infrared region indicative of circumstellar (CS) dust. Different models, among
which are a highly inclined passive flared disk with an inner rim (Dullemond et al. 2001); a flat
thick disk surrounded by a halo (Vinkovic et al. 2003); and a halo alone (Elia et al. 2004) have
been used to explain the spatial distribution of this dust. Notwithstanding their differences, all of
these models reproduced the observed SED reasonably well, which indicates the need for high-
resolution imaging to provide additional critical constraints.
Spatial observations at various wavelengths imply that the CS dust in the AB Aur system
lies in a disk and some type of more extended structure, Figure 4-1. An inhomogeneous envelope
extending to 1300 AU is apparent in optical scattered light (Grady et al. 1999), while closer to
the star, in optical and near-IR scattered light, one sees what appears to be a disk with quasi-
spiral structure, a radius of 580 AU, and an inclination of 300 (face-on = Oo), assuming flat
geometry (Grady et al. 2001, Fukagawa et al. 2004), Figure 4-2. CO observations reveal a
complex disk with an inner hole of about 70 AU extending out to 1000 AU and possibly having
non-keplerian motions, while 1.4 mm continuum observations indicate a disk with an inner
radius of 1 10 AU and an outer radius of 350 AU (Pietu et al. 2005; see also Mannings and
Sargent 1997; Corder et al. 2005). Near-IR interferometric studies resolve the inner 0.7 AU
region of the disk (Millan-Gabet et al. 2001). However, previous mid-infrared studies present
somewhat contradictory results. Marsh et al. (1995) report an extended structure at 17.9 Clm with
a semi-maj or axis of 80120 AU and an inclination of 750. Chen and Jura (2003), using the 10
meter Keck telescope, do not confirm that detection of extended structure, and at 20.5 Clm using
deconvolved images with a resolution of 0".6, Pantin et al. (2005) report an elliptical ring-like
structure at an average distance of 280 AU from the star. In addition, Liu et al. (2005) resolve the
inner disk interferometrically at 10.3 Clm, determining a size of 27 & 3 AU and an inclination of
450. There is also recent evidence that AB Aur could be the brighter component of a binary
system, with a companion separation most likely between 1 and 3 arcseconds (Baines et al.
In this Chapter I present deep mid-infrared images of AB Aur obtained at Gemini North. I
have resolved the emission close to the star at 11.7 and 18.1 Clm, and found an additional
extended component that appears to be roughly circularly symmetric. I show how these
observations of the thermal emission from dust in the AB Aur system help establish a more
coherent picture of the dust geometry consistent with most observations at other wavelengths.
Observations and Data Reduction
Observations of AB Aur were obtained on 2003 November 7 using Michelle (Roche 2004)
at Gemini North as part of the imaging survey of Herbig Ae/Be stars. Images, Figure 4-3, were
taken using the Si-5 filter (11.7 Clm, ah = 1.1 Clm) and Qa filter (18.1 Clm, ah = 1.6 Clm). The
standard chop-and-nod technique was used to remove thermal background from the sky and
telescope. The total on-source time was 645 seconds for each filter. The nearby point-spread-
function (PSF) star PPM 94262 was observed before and after each observation of AB Aur. The
average values of the full-width at half-maximum (FWHM) intensity of the PSF star were 0".28
at 11.7 Clm and 0".44 at 18.1 Clm, comparable to the diffraction limits (h/D) of the telescope at
these wavelengths. A Moffat function was fitted to the radial profile of each source to derive the
value of the FWHM. All observations were made at air masses less than 1.2.
Using the standard stars Sirius and Vega for flux calibration and air mass correction, I
derive flux densities for AB Aur of 21 & 2 Jy at 11.7 Clm and 36 & 4 Jy at 18.1 Clm. These values
fall within the spread of previous measurements for AB Aur, Table 4. 1, which is thought to be
variable in the mid-infrared, and the reported error bars are typical for mid-infrared photometry.
In addition to the emission from dust, the measured fluxes contain mid-infrared radiation from
the stellar photosphere. I estimated the contribution from the photosphere by using the NextGen
models developed by Hauschildt et al.(1999), assuming a stellar effective temperature of 9500 K
(van der Ancker et al. 1998), log g = 4.0, and solar metallicity. I scaled the model using UBV
fluxes for AB Aur (Malfait et al. 1998). From this procedure, I estimate the photospheric flux
densities to be 0.054 Jy at 11.7 Clm and 0.021 Jy at 18.1 Clm, which are negligible compared to
the emission from dust.
We show in Figure 4-4 the contour levels for AB Aur and the PSF star, scaled so that peak
fluxes and the lowest contours are at the same flux level for both images. The lowest contour
represents the 3 o flux level (three times the background noise) for the AB Aur images. The far
right panels in the figure show the residuals after subtraction of the normalized PSF from AB
Aur. This residual corresponds to fainter emission extending out to 2 arcsec (280 AU) from the
star at 11.7 Clm and to 2.5 arcsec (350 AU) at 18.1 Clm. This later value equals the outer
boundary of the dust emission seen in the millimeter continuum by Pietu et al. (2005). I do not
see evidence in the images of a ring-like structure as proposed by Pantin et al. (2005).
At these wavelengths a fit to the contour levels of AB Aur between 0".2 and 2".0 in radius
indicates an approximately constant position angle of 800 a 110, with an inclination (under the
assumption of flat disk geometry), i=cos-1(Ay/Ax), of 290 & 1 1o at 1 1.7 Clm and 12o a 12o at 18. 1
Clm, where Ay and Ax are the semi-maj or and semi-minor axis of the fitted ellipse. The
uncertainties in the measurements arise from multiple fittings as I vary the radii of the ellipse.
This low, nearly face-on, inclination is consistent with recent optical and near-IR results (Grady
et al. 01, Fukagawa et al. 04).
We also clearly resolve a bright, compact inner emission region near AB Aur. Figure 4-5
shows the values of the FWHM for AB Aur and the PSF star. I clearly resolve a bright, compact
inner emission region near AB Aur. Since integration time can influence the FWHM of the
image due to effects such as guiding error, rotation of the pupil and seeing, I divided the data into
subsets (equal to the nod dwell time) of equal integration time for AB Aur and PSF star. By
comparing these subsets, I conclude that the difference of the FWHM for AB Aur and the PSF
star is evident and not an artifact of integration time. Quadratic subtraction of the average
FWHM values of the source and the PSF star gives source sizes for the strong compact emission
of 17 & 4 AU at 11.7 Clm and 22 & 5 AU at 18.1 Clm. Error bars in these measurements were
calculated using error propagation of the standard deviation of the average FWHM values for AB
Aur and PSF star as seen in Fig. 4-5. The size of the compact emission is consistent with that
determined by Liu et al. (2005) at 10 Clm.
Even elementary considerations imply that the AB Aur disk is structurally complex. In
this section I show that radial variations probably exist in the grain size and/or composition, and
that the extended mid-IR originates in an optically thin region bounding an optically thick layer
near the disk mid-plane.
Particle Temperature and Sizes
Because the most extended mid-IR emission is faint, I divide the images into three annuli
centered on the star, each with a width Ar =100 AU. This is equivalent to 2.3 and 1.6 of the
resolution elements at 11.7 and 18.1 Clm, respectively. By assuming that the mid-IR emission is
optically thin and that the measured fluxes at both wavelengths originate within the same region,
I then calculate average color temperatures of 21513 K, 18916 K, and 184110 K for the
circumstellar dust within each of the three regions, with the highest temperature corresponding to
the region nearest the star. The quoted uncertainties in these temperatures are due only to
measurement errors in the flux densities, which are the relevant uncertainties when examining
the radial trends in temperature and corresponding dust properties. The expected blackbody
temperatures at these distances are 103 K at 50 AU, 59 K at 150 AU, and 46 K at 250 AU, which
are much lower than the derived color temperatures, indicating the presence of smaller, less
efficiently emitting grains.
The temperature of a dust grain in thermal equilibrium with stellar radiation depends on
the distance to the heating source and on the radiative efficiency of the dust particle. The value of
the efficiency depends on the properties of the material and the size of the grain and can be
considered equal to unity for radiation shorter than a critical wavelength ho. Given some
assumptions about the material, the value of ho can be used to estimate a characteristic grain
radius a as explained in detail in Chapter 3. One finds that ho a in the case of moderately
absorbing dielectrics such as graphite and amorphous silicate, while ho/2x: = a for strongly
absorbing dielectrics like polycyclic aromatic hydrocarbons (PAHs) (Backman and Paresce
1993). I used the energy balance equations as presented in Chapter 3 to calculate the expected
flux ratio (1 1.7/18.1) for different values of ho as a function of distance to the star and constrain
the sizes of the mid-IR emitting particles throughout the disk. From the observations, the
observed flux ratios for the three regions are 0.63-0.68, 0.46-0.53, and 0.44-0.51, with the values
decreasing as I move further from the star. Assigning these average flux ratio values to the
midpoint of each region and assuming moderately absorbing materials, I see that the emission in
these regions can be well constrained to come from particles of sizes 1.1-1.3 Clm in the central
100 AU, 0.2-0.3 Clm for the dust between 100 and 200 AU, and 0.08-0.12 Clm for the dust
between 200 and 300 AU. For strongly absorbing material, the corresponding values are 0.2-0.8
Clm, 0.02-0.05 Clm, and 0.008-0.02 Clm for the three regions respectively, Figure 4-6. For strongly
absorbing materials, the corresponding average values are 0.2 Clm for the compact emission, and
0.04 and 0.01 Clm for the other two regions.
These considerations suggest that I need different dust populations to fit the color
temperatures at different distances from the star. The circumstellar dust in AB Aur is probably a
combination of different materials, since PAH emission bands at 3.4, 6.2, 7.7, 8.6, and 11.3 Clm
and the amorphous silicate emission feature at 9.7 Clm have been detected (Cohen 1980, van den
Ancker et al. 2000). Thus, the deduced radial variation in particle size may result partly or
entirely from a radial variation in dust composition. Pietu et al. (2005) find that the CO emission
in AB Aur extends out to 1000 AU, whereas the millimeter dust continuum, such as the 18.1 mm
emission, extends only to 350 AU. They propose that this marked difference in the CO and dust
continuum radial distributions may be due to a fairly abrupt change in the radial variation in the
dust opacity, perhaps associated with less-processed dust at larger radii. Thus, at least
qualitatively, both sets of observations support the idea that there are radial variations in the dust
properties in the AB Aur disk. Mid-IR spectroscopy from the Very Large Telescope
Interferometer (VLTI) of three other Herbig Ae stars has revealed radial variations of dust
composition in those systems, with crystalline silicates dominating the inner 2 AU regions, and a
mixture of crystalline and amorphous silicates located in the 2-20 AU regions of the disks (van
Boekel et al. 2004). Detailed follow-up mid-infrared spectroscopy with high spatial resolution of
the AB Aur system can help resolve this issue.
Dust Optical Depth and Disk Morphology
Stellar optical radiation penetrates the CS dust to a radius corresponding roughly to zv = 1.
I assume that the disk is flat, with an inner radius R,,, = 0.5 AU, an outer radius Ro,,, = 400 AU, a
thickness AS = 10 AU, a dust mass M~,,,, = 2 x 1029 g (Mannings and Sargent 1997), and a
uniform volume dust density given by Equation 3.16.
We estimate from these considerations that zv = 1 at R 2 AU. The value obtained is an
upper limit for several reasons. First, for simplicity I have assumed a uniform volume dust
density; however, a radially decreasing density distribution with power-law indices in the range
0.5-2.0 is more consistent with previous models (Beckwith et al. 1990; Menshchikov and
Henning 1997; Dullemond et al. 2001) and will increase the dust density closer to the star.
Second, extrapolating values of the visible absorption coefficient from the submillimeter region
underestimates the visible extinction along the line of sight to the star, which is five times larger
than predicted by this relationship at 0.55 Clm (Mathis 1990); therefore, the value of zv =1 is
probably reached at R < 2 AU. I conclude that the stellar radiation does not penetrate very far in
the plane of the disk, in agreement with the conclusion of Mannings & Sargent (1997).
However, for the dust far from the star to reach the temperatures inferred from the observations,
it must be heated by direct radiation from the star, which implies heating of the dust above and
below the mid-plane. This dust could reside either in the surface layer of a flaring disk or in an
To address this issue further, I consider the vertical (i.e. perpendicular to the plane of the
disk) optical depth in the system. Since I believe the disk to be almost face-on, an estimate of
the vertical optical depth of the disk at mid-infrared wavelengths should indicate whether I am
only detecting emission from the dust located above the disk (if the disk is optically thick to mid-
infrared radiation) or from both sides of the disk (if the disk is optically thin to mid-infrared
radiation). The dust in the mid-plane evidently would be too cool for us to detect, since visual
radiation cannot penetrate there. Using the best fit to the AB Aur data for the disk surface
density distribution, Equation 3.19, where Co = 104 g Cm-2 (Dullemond et al. 2001) and r, = 7c,C,
I find that the AB Aur disk is optically thick vertically to mid-IR radiation out to a radius of 118
AU. The disk becomes optically thin to mid-IR radiation beyond 118 AU from the star. This
result suggests that, except in the outermost disk, I cannot look through the disk at mid-IR
wavelengths and are only detecting mid-IR emission from the surface" layer of the disk facing
us. This result has its parallel in the CO observations of AB Aur by Pietu et al. (2005), which
show that the optically thick emission lines, arising near the more directly irradiated disk surface,
have higher excitation temperatures than the optically thin ones weighted toward material in the
The passive flared disk model with inner rim developed by Dullemond et al. (2001) returns
the detailed geometry of a CS disk for given specific stellar parameters and general disk
properties such as dust mass and inner and outer disk radius. I have applied their model code
(kindly provided by C. P. Dullemond) to AB Aur to compare the inferred geometry of the disk
model to parameters derived from the observations. In this context, the size of the compact mid-
IR emission detected in the images coincides with the emergence at a radius of ~ 10 AU of the
disk from the shadow of the inner rim. The model predicts a jump in the surface temperature of
the dust at this boundary, which translates into stronger fluxes at mid-IR wavelengths. The dust
temperature implied by the images for the compact mid-IR emission is about 100 K lower than
the value predicted by the model for the surface layer of the disk at the onset of flaring (see
Figure 4-7). This inconsistency results from the fact that I am comparing an average color
temperature within the inner 100 AU to the peak temperature from the model within this region.
With the exception of the dust at the inner rim, which is heated to dust sublimation temperatures,
the model dust temperatures in the inner 10 AU are colder than 200 K because that dust is
shadowed from stellar radiation. The dust temperature peaks at a radius of about 10 AU and
decreases outwards as r-2. The model indicates that the average temperature of the surface layer
in this region is about 200 K, which is consistent with the results in the previous section. In this
context, I suggest that the strong compact emission detected in the images is a combination of
unresolved emission from the inner rim and emission from the surface layer of the disk at the
onset of flaring, which is resolved. Likewise, the size of the more extended component that I
detect in the images is as predicted from the model for the outer surface layer of the disk, which
becomes too cold at very large radii to be detected in the mid-IR.
The segregation of particle sizes that I derive from the observations can also be explained
within the passive flared disk model. Small grains in a region optically thin to stellar radiation,
like the surface layer of the disk, experience the effect of radiation pressure. The parameter P is
the ratio of the radiative to gravitational forces, which is proportional to the luminosity of the star
and inversely proportional to the particle size. For f > 0.5, particles are on unbound orbits and
will be expelled from the system in timescales comparable to the orbital period of the region
where the particles were produced (Backman & Paresce 1993). In the case of AB Aur, these
particles will be expelled in timescales of less than 103 years for a 500 AU disk in a gas-free
environment. I do not consider the presence of molecular CO at radial distances comparable to
the mid-infrared emission to be relevant for this result because I believe that the gas and the mid-
infrared emitting particles occupy different vertical regions of the disk. When gas molecules
collide frequently with dust particles, as is the case when abundant gas and dust coexist, they
reach similar temperatures; gas temperatures can drop below dust temperatures as gas is depleted
in the system. Pietu et al. (2005) derived temperatures of 70 K for the warm gas near the surface
of the disk at a radius of 100 AU and found no evidence for CO depletion. The derived dust
temperatures for this radial distance are substantially higher; therefore, it is appropriate to
assume that the mid-infrared emitting dust resides in an optically thin layer at the surface of the
disk, while the warm gas detected in the system lies at a lower elevation above the midplane.
Even if gas and dust coexist, Klar and Lin (2000) showed that hydrodynamic drag forces in a
gas-rich environment tend to enhance the effect of radiation pressure, increasing the critical size
for particles blown out of the system in a dynamical timescale. The presence of micron-size
particles in a system 2-4 million year old implies that there is a replenishing mechanism for these
grains, probably collisions of larger particles (Wyatt et al. 1999). Regardless of the formation
process involved, if small particles are constantly created, then I should see them at those
locations where they are produced or when they are driven outward by radiation pressure from
regions closer to the star. They should have lower emission efficiencies, and therefore, elevated
temperatures compared to larger dust at the same distance from the star. These smaller, hotter
particles would be easier to detect in the mid-infrared farther from the star.
I also considered the implications of the observations if the detected emission originated in
a disk and a halo, as proposed by Vinkovic et al. (2003). The presence of a halo could explain the
difference between the derived inclination angles from the study and that of Liu et al. (2005).
Given that the images achieved higher sensitivity than those of Liu et al. (2005), I could be
seeing more of the tenuous halo or spherical component, which will dilute the asymmetries
created by an inclined disk. However, the variation of grain sizes derived from the observations
is inconsistent with emission from a halo. In the roughest approximation, if a halo were present I
would expect to find a population of larger colder grains concentrated in the disk, and smaller
micron and sub-micron size particles residing in the halo. If this were the case, the average grain
size in the inner region should be smaller than at larger distances because I would be intercepting
more of the halo material in the central region, which is the opposite of the findings.
Our mid-IR images reveal two different emission components in AB Aur. The central
stronger emission is resolved, with quadratically deconvolved FWHM sizes of 17 & 4 AU and 22
& 5 AU at 11.7 Clm and 18.1 Clm, respectively. I also detect fainter extended emission out to a
radius of 280 AU at 11.7 Clm and 350 AU at 18.1 Clm. Emission is slightly elongated at 11.7 Clm,
indicating a disk inclination angle in the range 290 & 11o and a PA of 800 a 110. The morphology
at 18.1 Clm is more circular, corresponding to an inclination angle of 12o a 12o. However, within
the uncertainties, inclination angles at 11.7 and 18.1 Clm are the same. Assuming moderately
absorbing material, I derive average radii of the mid-infrared emitting dust in the system and find
that larger particles (a ~ 1 Clm) dominate the mid-infrared emission in the inner (<100 AU)
regions of the disk, and smaller particles (a < 0.3 Clm) dominate in the outer regions of the disk.
These results are reasonably well accounted for by a model of a passive flared disk with an inner
rim. The presence of a more spherical component fails to account for the particle size
segregation derived from the observations.
Table 4.1. Comparison of flux density measurements for AB Aurigae.
Date Observed Wavelength [Cpm] Flux [Jy] Reference
Oct. 1993 11.70 20.60 Marsh et al 95
Nov. 1994 11.70 23.00 Marsh et al 95
Feb. 2000 11.70 25.00 Chen & Jura 03
Aug. 2000 11.70 19.00 Chen & Jura 03
Nov. 2003 11.70 19.67 this study
Nov. 1994 17.90 34.70 Marsh et al 95
Feb. 2000 18.70 31.00 Chen & Jura 03
Aug. 2000 18.70 17.00 Chen & Jura 03
Nov. 2003 18.10 35.50 this study
Figure 4-1. Near-infrared scattered light of AB Aurigae taken with STIS at the Hubble Space
Telescope (Grady et al. 1999). The image shows the dusty disk closer to the star as
well as more extended nebulosity.
Figure 4-2. Near-infrared scattered light image of AB Aurigae showing spiral structure
(Fukagawa et al. 2004)
-4 -2 O
-4 -2 O2 4
Figure 4-3. AB Auriga false color images at 11.7 Cpm and 18.1 Cpm (left and right) taken with
Michelle on Gemini North. These maps clearly show the extended mid-IR emission.
The lowest color display is at the 30 level of the background.
IJ1Rur-lldrm C PSF-;a.a rrS! ~ :~.. b R~ldwrlr-l.l.cj~m
r :i.'L~,B~ ~4
Ir' r~ 5
"b' P: : ~
~s~La-IS,5rrm ~ t P5F-18,5~1m t ~R~sidr~al~-IB,~n;' ~p
rt ? i
$ I ss~.,
I ~ 1C r:c;
(~n s 'b;~~: a
+ + d. II,
-4 -1 D Z 4 iZ -2 ~ 2
4 -4 -Z o 1 4
Figure 4-4. Normalized contour level of AB Aur, PSF star and PSF subtracted emission. Upper
panels show 1 1.7 Clm data (contour) of AB Aur, PSF star scaled at 100 % of AB Aur
peak emission, and PSF subtracted emission. AB Aur and PSF star contour levels are
(0.06, 0.25, 1.07, 4.50, 18.86, 79.20) Jy/arcsec2, lOwest contour is 3 o (60
mJy/arcsec2) Of the background for AB Aur. Contour levels for the residual emission
at 11.7 Cim are (0.37, 0.69, 1.31, 2.47, 4.67, 8.83) Jy/arcsec2. Lower panels show 18.1
Clm data (contour) of AB Aur, PSF star scaled at 100 % of AB Aur peak emission
(Sirius was used to obtain better signal to noise ratio at this wavelength), and PSF
subtracted emission. Contour levels are (0.15, 0.33, 0.73, 1.61, 3.54, 7.77) Jy/arcsec2.
Lowest contour is 3 o (150 mJy/arcsec2) Of the background for AB Aur. Emission
from AB Aur can be seen extending to ~ 2" in the 11.7 Clm data and to ~ 2.5" in the
18.1 Clm data. For cosmetic reasons, a very noisy part of the image at the extreme
right of the 1 1.7 Clm PSF image has been removed; that has no effect in the region of
the stellar image, per se