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CO in the galactic center

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Title:
CO in the galactic center a complete survey of carbon monoxide emission in the inner 4 KPC of the galaxy
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Bitrán Carreño, Mauricio Ernesto, 1954-
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[s.n.]
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English
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x, 235 leaves : ill. ; 28 cm.

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Subjects / Keywords:
Galactic Center ( jstor )
Galaxies ( jstor )
Galaxy rotation curves ( jstor )
Kinematics ( jstor )
Latitude ( jstor )
Longitude ( jstor )
Milky Way Galaxy ( jstor )
Molecular clouds ( jstor )
Radial velocity ( jstor )
Velocity ( jstor )
Astronomy thesis Ph.D
Dissertations, Academic -- Astronomy -- UF
Galactic center ( lcsh )
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bibliography ( marcgt )
non-fiction ( marcgt )

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Thesis:
Thesis (Ph. D.)--University of Florida, 1987.
Bibliography:
Includes bibliographical references.
General Note:
Typescript.
General Note:
Vita.
Statement of Responsibility:
by Mauricio Ernesto Bitran Carreño.

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--~- ~{'


CO IN THE GALACTIC CENTER:
A COMPLETE SURVEY OF CARBON MONOXIDE EMISSION
IN THE INNER 4 KPC OF THE GALAXY






BY


MAURICIO ERNESTO BITRAN CARRENO
L -


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


1987

































Copyright 1987

by

Mauricio Ernesto Bitran Carreno
































This dissertation is dedicated with much love to my mother

Cora C. de Bitran, to my wife Gloria, to my son Jonathan, and specially to the blessed memory of my father Raul Bitran Nachary (Z.L.).















ACKNOWLEDGMENTS


This work was possible thanks to the efforts and

dedication of a number of persons and institutions. To all of them I express my sincere appreciation.

I thank Dr. Stephen Gottesman, my advisor at the University of Florida, for his excellent guidance, continuous support, and friendship during the course of this work. Likewise I thank Dr. Patrick Thaddeus, my coadvisor,

for his stimulating scientific guidance, hospitality, and unrestricted access to the computing facilities of the Goddard Institute for Space Studies. The supervisory committee members, Drs. T. D. Carr, J. Hunter, J. Ipser, A. G. Smith, H. Smith, and W. Weltner, are gratefully thanked for their useful suggestions and comments.

The project of millimeter-wave observations of the Southern skies with a one-meter class telescope was conceived and planned by Dr. Patrick Thaddeus. Its successful implementation was largely due to the efforts of

the project's manager Dr. Richard Cohen, with the help of Drs. Leonardo Bronfman, and David Grabelsky. J. Montani, F. Aviles, and M. Koprucu, with dedication and competence, kept the telescope running smoothly during its operation. Profs. Jorge May and Claudio Anguita, of the University of Chile,


iv









promoted and supported the involvement of Chilean astronomers in the project.

I am indebted to Dr. Leonardo Bronfman for advice, as well as for his hospitality and friendship. Many thanks go

to M. Koprucu and Dr. H. Alvarez for help with the observations, and to A. Smith, Dr. R. S. Cohen, Dr. L. Erickson, and Dr. T. Dame for help and advice with the data reduction.

I thank the University of Chile for support during the

observations at Cerro Tololo, and the faculty of the University of Florida for support and encouragement during my graduate studies. During this period we have enjoyed the friendship and support of Drew and Jill Weisenberger, Celia

Gottesman, and of fellow graduate students F. Reyes, G. Fitzgibbons, J. Webb, M. J. Taylor, and W. Cooke, in particular.

My deepest gratitude goes to my parents and to my wife for their permanent love and support. Dr. Gloria Rachamin,

my wife, deserves special recognition for her unfailing love, support, and encouragement, as well as for her skillful editing of the manuscript.


v

















TABLE OF CONTENTS


PAGE

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

ABSTRACT . . . . . . . . . . . . . . . . . . . . viii

CHAPTERS

I. INTRODUCTION . . . . . . . . . . . . . . . . . . 1

The Center of the Galaxy . . . . . . . . . 1 CO Observations of the Galactic Center . . 7 Present Work . . . . . . . . . . . . . . . 8

II. THE SURVEY . . . . . . . . . . . . . . . . . . 11

Instrumentation . . . . . . . . . . . . . 11
Observations . . . . . . . . . . *. .... 14
Large-scale Characteristics of the
CO Emission . . . . . . . . . . . . 17

III. WIDE-LINE CLOUDS NEAR THE GALACTIC CENTER . . 41

Kinematic Features or Molecular Clouds? . 43 The Masses . . . . . . . . . . . . . . . . 45
Origin of the Energy . . . . . . . . . . . 48
Are the Wide-Line Clouds Powered by
the Galactic Nucleus? . . . . . . . 52

IV. A CATALOGUE OF MOLECULAR CLOUDS IN THE
INNER GALAXY . . . . . . . . . . . . . 64

V. THE KINEMATICS OF THE MOLECULAR GAS IN THE
INNER GALAXY . . . . . . . . . . . . . 78

Introduction . . . . . . . . . . . . . . . 78
The Molecular Disk . . .. *.. . ...... 82
The Tilt of the Molecular Disk . . .*. 83 The Thickness of the Molecular Disk . . 92
The Rotation Curve of the Inner Galaxy . . 96


vi











VI. SUMMARY . . . . . . . . . . . . . . . . . . 120

APPENDICES

A. DETERMINATION OF MOLECULAR MASS IN CLOUDS . . 124

CO Masses . . . . . . . . . . . . . . . 124
LTE Masses . . . . . . . . . . . . . . . 126
Virial Masses . . . . . . . . . . . . . 128
Molecular Masses near the Galactic Center 129 B. LONGITUDE-LATITUDE MAPS . . . . . . . . . . . 131

C. LONGITUDE-VELOCITY DIAGRAMS . . . . . . . . . 147

D. LATITUDE-VELOCITY DIAGRAMS . . . . . . . . . 177

REFERENCES . . . . . . . . . . . . . . . . . . . 229

BIOGRAPHICAL SKETCH . . . . . . . . . . . . . . . 235


vii















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



CO IN THE GALACTIC CENTER:
A COMPLETE SURVEY OF CARBON MONOXIDE EMISSION
IN THE INNER 4 KPC OF THE GALAXY


By


Mauricio Ernesto Bitran Carrefio


December 1987


Chairman: Stephen T. Gottesman Cochairman: Patrick Thaddeus Major Department: Astronomy



The first well-sampled, large-scale survey of 12CO (J=1-0) emission from the inner 4 kpc of the Galaxy is presented and used to study the distribution of molecular clouds and the kinematics of the molecular gas in the inner Galaxy.

The survey samples a 40 wide strip along the Galactic

equator from 1=-120 to 1=130. The over 8000 spectra obtained with the Columbia University Southern millimeter telescope (La Serena, Chile) have a velocity resolution of

1.3 km s- , a rms sensitivity better than 0.12 K, and are


viii









spaced by approximately one beamwidth (8.8'). This is the first survey to encompass the complete latitude and velocity spans of the CO emission from the inner Galaxy.

The survey is presented as a collection of 1-V, b-V, and 1-b maps. Several new CO features were observed, and

the molecular counterparts of the classic H I structures appear with unprecedented clarity owing to the dense sampling, high sensitivity, and extended latitude coverage of this survey.

Features with the largest CO luminosities and velocity widths in the inner Galaxy, outside the nuclear region, were fully mapped and analyzed. The largest of these objects are located at (1, b) = (3.20, 0.30), and (5.30, -0.30), in a zone previously considered almost devoid of CO emission. Their respective CO luminosities exceed 8% and 3% of the CO

luminosity of the inner 500 pc of the Galaxy, integrated over all observed latitudes and velocities. Their velocity widths reach 140 and 100 km s- (FWHM), respectively.

We argue that these features are peculiar molecular clouds located in the vicinity of the Galactic center, and that their large internal kinetic energies (of the order of

1054 ergs) may have originated in a Seyfert-like event at the Galactic nucleus about one million years ago.

Molecular clouds in the surveyed area were identified

and catalogued. The inclination and thickness of the molecular layer in the inner Galaxy were measured and their


ix









relevance to current models of the region discussed. A lower limit for the surface density at the center was found,

and a CO rotation curve was calculated and compared to available H I rotation curves.


x















CHAPTER I
INTRODUCTION



The Center of the Galaxy

The Galactic center has aroused the interest of

astronomers ever since it became clear that the Earth and the Sun are far from the center of the Milky Way, and that a

massive galactic nucleus exists in the direction of the constellation of Sagittarius (Shapley 1916 a, b; 1919 a, b, c). The center, however, is hidden from our view by interposing clouds of dust and gas that absorb most of the visible and ultraviolet light emitted in our direction. As a consequence, most studies of the Galactic center have had

to rely on radio, infrared, X-ray, and y-ray observations.

In a paper presented in 1935, reporting radio

observations of the sky at 20 MHz, Jansky noted that "radiations are received every time the antenna is directed

towards some part of the Milky Way system, the greatest response being obtained when the antenna points towards the center of the system" (Jansky 1935, pg. 1158). This was the first observation of the Galactic center, which also marked the birth of radio astronomy.

Following this pioneering work (Jansky 1932, 1933,

1935), the Galactic center region has been observed at many


1







2


different wavelengths, with increasing sensitivity and resolution, in the continuum and in spectral lines. An enormous volume of data has accumulated. However, because

of the complexity of the region, the development of a physical picture that can coherently accommodate all this information has been slow (see reviews by Oort (1977) and Brown and Liszt (1984)).

The Galactic center is a singular region. It displays

intense activity and characteristics unlike those observed anywhere else in the Galaxy. At the center, near-infrared observations indicate the presence of a spheroidal

distribution of late-type stars, in which the density increases inwards (Allen et al. 1983), reaching a maximum near the infrared source IRS 16. Coinciding with the central star cluster and extending some 80 light years in diameter lies Sgr A, a highly complex radio source which is very intense in both thermal and nonthermal emission. Sgr A is also very prominent in molecular line and X-ray emission

(Watson et al. 1981), and may be the source of a variable y-ray line produced by positron-electron annihilation at 511 KeV (Lingenfelter and Ramaty 1982).

High resolution observations at centimeter wavelengths (Ekers et al. 1983) resolve Sgr A into three components: Sgr

A East, Sgr A West, and Sgr A*. The shape and spectral index of Sgr A East indicate that it is probably a supernova remnant. Sgr A West is a strong source of thermal







3


radiation in which filaments of high-density ionized gas move at radial velocities in excess of 200 km s-1. If these

motions are due to gravitation, the observed speeds would indicate a mass of 3x106 Mo within 0.5 pc of the center (Lacy et al. 1980). At the centroid of Sgr A West lies Sgr A a nonthermal compact radiosource with a diameter of less

than 20 A.U., a spectral index of -0.25, and a luminosity of 1034 ergs s-1 (Lo et al. 1985). Its emission

characteristics are best explained, according to Lo (1986), by low level accretion onto a massive (-106Mo) Black Hole. A Black Hole at the Galactic nucleus was proposed by LyndenBell and Rees (1971) to be the source of energetic phenomena in the central region of the Galaxy.

At a larger scale, the most striking feature is the radio continuum source known as the "arc" (Yusef-Zadeh, Morris, and Chance 1984). Formed by several straight parallel filaments about 100 ly long and 3 ly wide, it is oriented perpendicular to the'Galactic plane and apparently connected to Sgr A by curved filaments (see Fig. 1, YusefZadeh et al. 1984). The "arc," which has a spectrum with thermal and nonthermal components (Yusef-Zadeh et al. 1986), may be related to large scale magnetic fields in the center region.

Within a few hundred parsecs of the center, there is a

large concentration of massive molecular clouds with high velocity dispersions. Observations of the hydroxyl radical







4

(OH) (McGee et al. 1970, Cohen and Dent 1983), formaldehyde (H2CO) (Scoville, Solomon, and Thaddeus 1972; Cohen and Few

1981), ammonia (NH3) (Kaifu et al. 1975), carbon monoxide

(CO) (Bania 1980, 1977; Liszt and Burton 1978; Heiligman 1982), and other molecular tracers, indicate that these dense clouds are rich in molecular material; their main component presumably being molecular hydrogen (H2). The distribution of the molecular gas is strongly asymmetrical in both longitude and velocity; most of the emission occurs at positive velocities and longitudes. Some of the clouds in this region apparently form a rotating and expanding ring

(Kaifu, Kato, Iguchi 1972) with a radius of 190 pc, expanding at an average velocity of 150 km s-1 and rotating with a velocity of 65 km s~ (Bania 1980). Superimposed on this emission, a relatively symmetrical structure is

observed in H I and in CO. This feature has two highvelocity wings of emission located in opposite quadrants about the center (see Figs. 11-8 and 11-9), which were interpreted as the signature of a gaseous rotating "nuclear disk" (Rougoor and Oort 1960).

The molecular counterpart of Sgr A seems to be a very intense, overlying molecular cloud with a mean velocity of

40 km s-1. This "+40 km s-1 cloud" contains the most intense CO lines in the central region and appears to be connected by a bridge of CO emission to another remarkable molecular cloud of the central region, the Sgr B2 complex.







5

This unique object contains several compact H II regions and

is extremely rich in complex molecules, owing to the combination of large mass with high density. Most of the known interstellar molecules were first observed here.

Peculiar phenomena in the central region are found even at a larger scale. Vast gaseous features are seen engaged in large noncircular motions up to 4 Kpc from the Galactic

Center (Oort 1977, Rougoor and Oort 1960). The most

conspicuous of these is the "3 Kpc arm;" a large structure spanning galactic longitudes from about +120 to -220, where it apparently reaches a tangential point. It is observed in absorption in the direction of the Galactic nucleus at -52

km s-1 in both H I (Van Worden et al. 1957) and CO (Bania 1980). This feature is thought to be a spiral arm located some 3 kpc in front of the nucleus, and expanding away from it. Another large feature extending across 1=00 with a noncircular velocity is the "+135 km s-1 arm." It is probably

located behind the nucleus, since no absorption has been observed, and it also seems to move away from the center (see Figs. 11-8 and 11-9). Together, the "3 Kpc arm" and the "+135 km s-1 arm," constitute the largest deviation from

circular motion in the Galaxy, each containing a mass of order 2 x 107 Mo, and a total expansion energy of 4 x 1054 ergs (Bania 1980).

The interpretation of these large, apparently

expanding, gaseous structures is quite complex. Attempts to







6


explain them include expulsion from the nucleus (Van der Kruit 1970, Cohen and Davies 1976), a rotating and expanding

tilted disk (Liszt and Burton 1978), elliptical streaming motions (Peters 1975), gas in dispersion orbits (Simonson and Mader 1973), an eccentric bar-like structure (Burton and Liszt 1980), an oval distortion of the central region of the Galaxy (Yuan 1985), and gas moving in a triaxial potential

(Heiligman 1982, Vietri 1986, Mulder and Liem 1986). In each case there are features not satisfactorily explained by the models, and the underlying dynamical mechanisms in some of them are not clear. Thirty years after their discovery,

the "expanding features" remain some of the most puzzling characteristics of the central region of the Galaxy.

In summary, the Galactic center is "sui generis," a region with unusual behaviour observed in scales ranging from a few astronomical units to several kiloparsecs. It displays some characteristics similar to those of active galactic nuclei, although at a much lower level. The study

of the Galactic center is not only fundamental to the understanding of the structure and dynamics of our own Galaxy, but may also provide important clues on the phenomenon of nuclear activity in spiral galaxies.







7


CO Observations of the Galactic Center

Observations of the interstellar material in the central region of the Galaxy are the main source of information on the mass density and the physical conditions

prevalent at the center; an important fraction of our knowledge about this region has come from observations of interstellar molecules.

Although H2 is the most abundant interstellar molecule, it is very hard to observe directly from the ground. H2 is an homoatomic molecule, therefore, it lacks a dipole moment and dipole transitions. Near IR quadrupole rotationvibrational transitions can be observed from the ground, but this radiation arises only in quite hot regions. The

easiest way to study the cold molecular gas in the Galaxy is

to observe the dipole rotational transitions of the second most abundant molecule in molecular clouds, CO. Collisions with H2 molecules readily excite rotational states of the CO molecules. The first rotational state of the CO molecule, corresponding to an excitation temperature Tx = 5.5 K, is thermalized (Tx = Tk) by collisions at H2 densities of a few hundreds per cm3. The fundamental transition J=1-,0 is readily observed at 2.6 mm (115.2 GHz). This transition is

an acknowledged primary tracer of H2 in molecular clouds (Thaddeus 1977, Dame 1984, Scoville and Sanders 1986).

Most of the CO observations of the Galactic center to

date were obtained from the Northern Hemisphere with







8

instruments of relatively high angular resolution (-1' ). As a result of these constraints, observations of the latitude

distribution of molecular gas within a few kiloparsecs of the Galactic center are moderately to severely undersampled. The better sampled observations were restricted to small areas around the nucleus and near the Galactic plane

(Inatani 1982; McCutcheon, Robinson, and Whiteoak 1981; Heiligman 1982). At a larger scale, however, the out-ofplane observations consisted of longitude strips at a few latitudes, with observations typically spaced by 0.50 (~30 beamwidths) (Liszt and Burton 1978; Bania 1980; Sanders, Solomon, and Scoville 1984, Bania 1986). Grids this coarse can easily miss molecular clouds located near the center.

The results of the out-of-plane survey of the first galactic quadrant by Cohen et al. (1980), emphasize that well-sampled observations (sampling interval -1 beamwidth) are necessary in order to detect the continuity of largescale features. These considerations suggest that new, important information on the Galactic center region can be

gained from a large scale, well-sampled survey of the latitude distribution of CO emission.





Present Work

On the premise that out-of-plane features are allimportant for the understanding of the kinematics of the







9

central region of the Galaxy and impose critical constraints on models, the main goal of this dissertation is to make a

sensitive, homogeneous, and well-sampled survey of the latitude distribution of CO emission in the inner 4 kpc of the Galaxy covering all significant out-of-plane and highvelocity emission.

For this purpose, we have used the Columbia Southern millimeter-wave telescope at Cerro Tololo Interamerican Observatory, in Chile, to map an area of 100 sq. deg. with an 8.8' beam. The observations are spaced by 0.85 beam within 10 of the Galactic plane, where most of the CO emission is concentrated, and by 1.7 beams for 10 < |bI 5 20.

The combined advantages of observing from the Southern Hemisphere and of using a telescope with a relatively large beam, have allowed us to significantly increase the sampling

density, improve the sensitivity and extend the spatial coverage with respect to previous observations.

The main questions that this work addresses are

1. What is the large scale, two-dimensional (l,b)

distribution of molecular gas in the inner 4 kpc of

our Galaxy?

2. What are the physical characteristics of the

molecular gas in the center region?

3. Are there any significant out-of-plane features

missed in previous observations?







10

4. Is the paucity of molecular gas observed outside the

nuclear disk partially caused by lack of latitude

information?

5. Do these observations impose new constraints on

current models of the center region?

This dissertation is organized in six chapters and four appendices. Chapter II describes the Columbia Telescope, the observational techniques used, and presents the large scale characteristics of the molecular gas. Out-of-plane features of great CO luminosities and velocity widths, fully sampled for the first time in this survey, are discussed in Chapter III. Chapter IV presents a catalogue of molecular clouds in the central region, and their identification with H II regions when possible. The kinematics of the molecular gas at the Galactic center are discussed in Chapter V. In

this chapter current models of the region and their compatibility with present data are also discussed. A general summary is presented in Chapter VI, and the survey maps are presented in the appendices.
















CHAPTER II
THE SURVEY



Instrumentation

The Galactic center CO survey was carried out with the

Columbia Southern Millimeter-wave Telescope (Cohen 1983; Bronfman et al. 1986), located at Cerro Tololo InterAmerican Observatory (La Serena, Chile). Figure II-1 is a

diagram of the telescope, and a summary of its main specifications is found in Table II-1.

This Cassegrain Telescope, a fairly close copy of the Columbia Millimeter-wave Telescope in New York City (Cohen

1977), is an ideal survey instrument. At 115 GHz, the frequency of the CO J=l-10 rotational transition, its 1.2 m aperture gives a beam 8.8' wide (FWHM). This beam size allows dense sampling of a large area of the sky in a practicable period of time, yet it is small enough to resolve a typical giant molecular cloud located over 20 kpc away (the linear resolution at the Galactic center (Ro=8.5 Kpc) is 18.5 pc).

Computer control of the altitude-azimuth mount and the dome allows fast pointing changes: switching between

positions five degrees apart takes less than one second. This allows the use of reference positions located several


11







12


degrees away from the Galactic Equator with little idle time.

The telescope is equipped with a sensitive and stable superheterodyne receiver which has a single-sideband noise temperature of 380 K. The receiver's first stage, consisting of a Schottky barrier diode mixer and a GaAs FET amplifier at the IF of 1390 MHz, is cooled to 77 K by liquid nitrogen, while the second stage of amplification is kept at room temperature. Owing to the low noise temperature of the system and the good atmospheric conditions at Cerro Tololo,

integration times are typically quite short; a rms noise temperature uncertainty of 0.1 K is generally reached in less than 10 minutes.

The spectrometer is a 256-channel filter bank of standard design. Each filter, 0.5 MHz wide, affords a velocity resolution of 1.3 km s-1 at 115 GHz. The total spectral range of the spectrometer is 333 km s-1, adequate to cover most of the Galactic CO emission, except within a

few degrees of the Galactic center, where the CO emission reaches velocity widths up to 500 km s-1. In this region, two spectra centered at different velocities were taken at

each position, and combined as explained in the next section.

The atmospheric conditions at Cerro Tololo are

extremely favorable for millimeter-wave observations. The water vapour opacity per air mass is usually less than 0.1,







13

allowing observations virtually year around. An added bonus in the case of this survey is the favourable position of the Galactic center as seen from Cerro Tololo. The center transits almost overhead; so most of the spectra were taken through an air mass of about unity.

The system was calibrated daily and checked against reference CO sources before the start of the observations.

A standard blackbody chopper-wheel technique (Kutner and Ulich 1981) was used to produce antenna temperatures, TA*' corrected for atmospheric attenuation, resistive losses, scattering, and rearward spillover. To produce radiation temperatures, TR, the antenna temperatures should be divided by the telescope's main beam efficiency, p=0.82. A complete account of the calibration procedure and results is given by Bronfman (1986).

The pointing of the telescope was checked twice a day by observing the sun. Right ascension and declination scans across the sun's limb were used to locate the center of the solar 115 GHz continuum emission disk to within 15". The long-term pointing accuracy was checked every few months by

observing star positions with a small optical telescope which is collimated with the radio telescope (see Fig. II-1). This procedure ensured that long-term pointing errors were smaller than 1' (for details see Grabelsky 1985).







14

A Nova 4/x minicomputer with 128 kilobytes of memory was used to control the positions of the mount and dome, the frequency of the local oscillator, the data acquisition, and to perform on-line data processing. The computer system and

telescope control are described in detail by Grabelsky (1985).





Observations

The almost 8000 spectra comprised in the present survey were taken between February and November of 1984. The area observed, a strip of 250 x 40 along the Galactic plane, was sampled every 7.5' (0.85 beam) in both Galactic coordinates

within 10 of the equator and every 15' (1.7 beam) beyond. Relevant parameters of the survey are listed in Table 11-2.

Position switching was used for all observations in

order to subtract instrumental and sky background. This was

done by pointing the telescope alternatively to a target ("on") position and to a CO-free reference ("off") position

every 15 seconds and then subtracting the spectrum of the "off" position from that of the target position. The reference positions, selected from obscuration-free regions

found in the ESO plates, were verified to be free of CO emission greater than 0.04 K, which is about one-third of the rms noise per channel in the final spectra. Table 11-3

lists the CO-free ("off") reference positions used in the







15

survey. Position switching against clean "offs" yields very flat base-lines; so only first order base-lines were subtracted from the spectra.

The velocity centroid of CO emission varies strongly with Galactic longitude, changing sign at 1=00. Full velocity coverage required setting the spectrometer to different central velocities in the appropriate longitude ranges (Table 11-2). Nevertheless, within a few degrees of the Galactic center, the velocity width 6f the CO emission exceeded the spectral range of the spectrometer. This problem was circumvented by taking two spectra, centered at velocities 306 km s~1 apart, at each position within 50 of the Galactic center. The two spectra cover a combined spectral range of 640 km s-1 with an overlap of 26 km s1 (20 channels); this allows ample emission-free sections to set baselines, and enough overlap to check that the spectra match properly. The values of antenna temperature in the overlapping channels were checked to coincide within the rms temperature uncertainty of the spectra. Poor matchings were rejected and reobserved. A combined spectrum obtained with the procedure described above is shown in Fig. 11-2. The high signal-to-noise ratio and flat baseline of this spectrum are characteristic of the whole survey, as can be deduced from the representative spectra shown in Fig. 11-3.

An important goal of this survey was to cover the full latitude extent of the Galactic CO emission. We found most







16

of the emission to be confined between b=-10 and b=i0, and

therefore our latitude coverage (-20 < b < 20) seems adequate. However, since H I emission is observed at higher latitudes in opposite Galactic quadrants (at b>00 for 1<00, and at b<00 for 1>00), a search was conducted between b=20 and b=2.50 for 1<00, and between b=-20 and b=-2.50 for 1>00. Only local emission (V-0 km s~ ) was detected.

The basic datum of the survey is the antenna

temperature observed from a particular direction in the sky and at a given radial velocity, T(li, b1, Vk). Twodimensional representations of this four-dimensional data array were accomplished by projecting the integral of the temperature over one of the independent variables onto the

plane of the other two. Contour diagrams representing integrated temperature in the longitude-latitude (l,b), longitude-velocity (1,V), and latitude-velocity (b,V) planes

were used to exhibit different characteristics of the CO emission (see appendices B, C, and D respectively). In the l,b diagrams, velocity-integrated temperature is displayed in the plane of the sky. This representation is useful to compare the CO distribution with that of other tracers that offer no velocity information such as the infrared or radiocontinuum. Kinematic characteristics of the molecular gas

are best displayed in 1,V diagrams of latitude-integrated temperature, and the out-of-plane gas distribution is best seen in b,V diagrams of longitude-integrated temperature.






17


Large-scale Characteristics of the CO Emission

This section describes the general characteristics and the main features of the CO emission observed in the present survey. The distribution of molecular hydrogen mass in the plane of the sky is outlined in maps of velocity-integrated CO antenna temperature (WCO = J T dV ) in Galactic coordinates (Fig. 11-4). Longitude profiles of CO emission (I(l) = f J T db dV ), integrated over different latitude ranges, are displayed in Figs. II-5, 6, 7, and the main kinematic features observed are identified in 1-V diagrams integrated over all the observed latitudes (Figs. 11-8, 9).

The spatial distribution of CO emission in the inner

3.6 kpc of the Galaxy is displayed in Fig. 11-4. To examine

the full dynamic range of the observations, 1-b maps with three different contour level values (approximately 4, 10, and 20 standard deviations, respectively) are shown in Fig. 11-4. In the three maps of Fig. 11-4, the same CO emission data were integrated over the full velocity range (-320 < V < 320 km s1 ). As can be seen in Fig. 11-4, the CO

emission is distributed closely along the Galactic plane with the highest intensity concentrated between 1=-1.50 and 1=1.80 and within approximately 0.60 of the Galactic plane. This central source encompasses the large molecular

complexes Sgr A and Sgr B, the nuclear disk, and the expanding molecular ring, which are more clearly demarcated in 1-V diagrams (Figs. 11-8, 9).







18

In longitude profiles (Fig. II-5, 6, 7) the strongly peaked central source appears very asymmetrical with most of

the emission found at positive longitudes, resulting in an intensity-weighted mean longitude of 1=0.40. This central region can be resolved into six prominent peaks (Fig. II-5), with the most intense maximum located at 1=1.20. The central source is also asymmetric in Galactic latitude with most of the CO emission originating below the IAU Galactic plane (see Figs. 11-6 and 11-7 integrated below and above the Galactic plane, respectively), yielding an intensityweighted mean latitude of b=-0.050.

The total H2 mass contained in the central region (1 x

b = 3.30 x 1.30 = 490 x 190 pc), calculated from the CO luminosity (see Appendix A) under the assumption that all the emitting gas is at the Galactic center (Ro= 8.5 kpc), is 2.7 x 108 Mo. This value is in agreement, within a factor of approximately 2, with recent determinations of H2 mass at the center (see Table 11-4). It should be noted, however, that all the masses quoted in Table 11-4 might have been overestimated by a factor of 10, since the conversion factor from CO luminosity to H2 column density in the center may be an order of magnitude smaller than that in the disk of the Galaxy (for details see Appendix A).

Most of the CO emission of the inner Galaxy is contained in a band approximately 20 wide in latitude, astride the Galactic plane, and extending through all







19


observed longitudes. The emission occurs preferentially above the plane for negative longitudes and below the plane for positive longitudes. This asymmetry, also observed in

H I, results from a tilt of the gas distribution with respect to the Galactic plane. The magnitude of the tilt, however, is controversial (see Chapter V).

The most intense CO emission of the inner Galaxy, outside the central source, originates at two features located at 1~30, b>O0 and l~50, b<00 (Fig. 11-4). In

longitude profiles (Figs. II-5, 6, 7) they can be seen as two large peaks at 1=3.20 and 1=5.30, which correspond to projected Galactocentric distances of 477 pc and 786 pc, respectively. These sources are remarkable, not only for their high intensities, but also for their unusually large velocity widths (Fig. 11-8). Their combined luminosity reaches over 1/10 of the intensity of the central source and they are seen in a zone previously considered "almost devoid of CO emission" (Bania 1986, pg. 873). The full extent and intensity of these features were missed in other CO surveys of this region owing to undersampling. The characteristics

and origin of these "wide-line" features are discussed extensively in Chapter III.

Other discernible peaks outside the central region observed in Fig. 11-6 at 1=-3.80, -4.40, and -5.30 correspond to smaller wide-line features (see Table III-1). The main peak in the central source at 1=1.20, b>O0







20

(Fig.II-7) appears remarkably similar in shape and intensity to the wide-line feature at 1=3.20. This similarity may suggest that this peak is another wide-line feature concealed by the complex emission of the central region.

The kinematics of the CO emission in the inner Galaxy are characterized by asymmetry and noncircular motions. The asymmetry in velocity can be clearly seen in Fig. 11-8, where most of the emission occurs at positive velocities with an intensity-weighted mean velocity of 23 km s-l for the central source. If all the molecular gas in the Galaxy

were engaged in circular differential rotation around the Galactic center, the CO emission displayed in the 1-V plane

would take place only in the quadrants where velocity and longitude have equal signs. However, large features are also seen in the other two quadrants (Fig. 11-8) indicating the presence of gas at noncircular, "forbidden," velocities. These deviations from circular motion were discovered in the first H I surveys of the center.

Figure 11-8 displays the kinematic arrangement of molecular gas in the inner-Galaxy in the 1-V plane, integrated over all observed latitudes. Owing to its dense sampling, high sensitivity, and extended latitude coverage,

the present survey has revealed new CO features and the molecular counterparts of classic H I features with unprecedented clarity. The following is a brief description of the CO features observed in the 1-V plane, in order of







21


their relative intensities. To aid in the identification, the main features observed in Fig. 11-8 have also been represented schematically in Fig. 11-9. The names of the H

I features used by Van der Kruit (1970) are given in parenthesis when appropiate.

1. The main maximum. An intense ridge of emission seen

along all observed longitudes at low radial

velocities (V -0 km s- ). Its large latitude extent

and low velocity (see Appendix C) suggest that the

emission originates mostly from local material.

2. The nuclear disk (VdK IV). A high-velocity wing of

emission that extends in the plane from 1=-1.50, V=-230 km s-l towards the center and reappears at positive longitudes and velocities reaching 1=1.50,

V=220 km s-1 (Fig. 11-8, 9). This high-velocity

feature is symmetric about the center and was

interpreted as a gaseous disk in circular rotation

about the center by Rougoor and Oort (1960).

3. The 3 kpc arm. A long ridge of emission extending

from 1=130, V-0 km s-l across the whole 1-V diagram to 1=-120, V=-130 km s~ (Figs. 11-8, 9). Crossing 1=00 the 3 kpc arm is seen in absorption against Sgr

A at V=-52 km s-1; therefore it is located in front of the nucleus and appears to expand away from it.

The 3 kpc arm is the brightest H I feature displaying

noncircular motion; it extends to 1=-220 where it







22

apparently reaches a tangent point (Cohen 1975). The

3 kpc arm has been modeled as a ring sector (Bania 1980) and as waves excited at the resonance by an

oval distortion of the central region (Yuan 1985). No

signs of star formation have been observed in it

(Lockman 1980). See, however, chapter IV.

4. The 135 km s71 expanding arm (VdK I). A large

structure extending from 1=130 to 1=-50 at positive

velocities. It crosses 1=00 at 135 km s~1 and no absorption is observed. Therefore, it is probably located beyond the nucleus and apparently expanding away from it. It has been modeled as a ring sector

by Bania (1980).

5. The molecular ring. This ring of molecular clouds at

the Galactic center (not to be confused with the

molecular ring at R= 4 to 6 kpc), is a large curved structure mostly located at negative velocities near

the center, partially superimposed on the nuclear

disk. It extends from 1=-1.20, V~0 km s~ to

1=-0.50, V-140 km s1 and curves back towards the V=0 km s1 at 1=1.50. The molecular ring extends below the plane to b=-0.50 (Figs. C-17,19). This

feature has been modeled as an expanding and rotating

molecular ring by Scoville et al. (1972) and Bania

(1980). The prominence of this feature in

formaldehyde absorption (Bieging et al. 1980)







23

indicates that it is probably located in front of the

nucleus.

6. Sgr A. This prominent molecular cloud at V-40 km

s-1 has been identified with the Galactic center

because of its location and high intensity. A deep

absorption feature may indicate that this cloud is

partially in front of Sgr A*.

7. Sgr B and B2. Correspond to the large molecular

complex at 1~0.50, V-60 km s-1. One of the most

intense sources observed contains several compact H

II regions, complex molecules, and is the site of

active star formation.

8. The central gas layer. An intense bridge of CO

emission, appearing to connect Sgr A and Sgr B along the Galactic plane at positive velocities (see Figs.

C-15, 17). This apparently continuous layer is a unique object in the Galaxy because of its smooth

space distribution and large velocity width.

9. The connecting arm (VdK IIIa). Rougoor (1964) called

this feature the connecting arm because it seems to connect the nuclear disk, at 1=2.50, V=270 km s-1 to the main maximum around 1=130. The connecting arm joins the nuclear disk below the plane at b=-0.250

and, as it approaches the main maximum, its velocity

and mean latitude diminish. It can be clearly

identified in the 1-V diagrams of Appendix C (Figs.







24

C-17 to 25) and in the 1-b maps of Appendix B (Figs.

B-1 to 5).

10. The 165 km s~i feature (VdK XV). Although discovered

in H I, this feature is more prominent in CO. It

crosses 1=00 at V-165 km s- in the plane; since no

absorption against the nucleus is observed, this

feature is probably located behind the nucleus

apparently moving away from it. It has been proposed

to form part of the molecular ring (Scoville et al.

1972).

11. Norma and Scutum arms? Probably correspond to the

features seen betwen the 3 kpc arm and the main

maximum, at negative longitudes. They appear similar to the 3 kpc arm in longitudinal extension and slope

in the 1-V plane. Although fairly confused, two structures can be identified in this region and

followed towards the center. Absorption features at

V=-30 km s-1 and V=-10 (Fig. D-28) indicate that

these structures probably lie in the near side of the

Galaxy. A string of clouds observed between l=-20 and 1=20 at V-15 km s1 at b=-0.380 (Fig. C-18) has been identified by Heiligman (1982) with the ScutumCrux arm. The string of clouds between l=-40 and

1=-120 and about V=-30 km s1 could be the

continuation of the arm, while the clouds in the same region at about V=60 km s1, may belong to the Norma

arm.







25

12. Van der Kruit XIV. A symmetric counterpart to the

connecting arm at negative longitudes and velocities.

It can be seen in Figs. C-12 to C-15 at 1=-20 to -50

and between V- -220 to -150 km s1 It joins the disk at b=00 and continues above the plane towards the main maximum. It is much weaker and less well defined than the connecting arm. However, this is

one of the few features displaying symmetry about the

center.

13. Van der Kruit II. Also known as the 70 km s1 arm,

is a faint and narrow feature in the plane. It has

been studied in detail by Shane (1972).

14. Feature IIIb. An in-the-plane feature which runs

parallel to the 70 km s-1 arm and seems to end near

the wide-line feature at 1=5.30.

15. Van der Kruit IX. Found at higher velocities than

feature IIIb, crosses the connecting arm and becomes confused with emission of near-by features near 1=70 16. Van der Kruit XII. A weak cloud detected at 1=40,

b=-2.40, and forbidden velocity V=-83 km s~1, seems to be the first CO detection of feature VdK XII, the

most prominent high-velocity feature seen in H I

(Cohen 1975). A nearby cloud at V-30 km s1

coincides with the 1-V locus of the 3 kpc arm but

lies well below the plane ( b=-2.20). Not shown in

Fig. 11-9.







26

17. A feature extending below the plane between 1=3.30

and 4.00 and between velocities V=40 and 150 km s-l was detected for the first time in this survey. It could be related to H I feature J4 (Cohen 1975) but

their slopes in the 1-V plane are different (Fig.

D-19). Not shown in Fig. 11-9.

The characteristics of the molecular clouds that form

many of these features are tabulated in a catalogue in Chapter IV.







27


TABLE 2.1


TELESCOPE CHARACTERISTICS


Aperture Beam width (FWHM) Effective F/D Beam Efficiency Noise Temperature Spectrometer


1.2 m

8.8' at 115 GHz

3.79

0.82

380 K Single sideband 256 channels 0.5 MHz wide







28


TABLE 2.2


PARAMETERS OF THE SURVEY







Angular resolution.....8.8' Velocity resolution.... 1.3 km/s Sensitivity............0.12 K rms per velocity element


Spatial coverage:

Galactic longitude.....3480 to 130 Galactic latitude......-2.00 to 2.00 Sampling interval:

0.00 < [b] < 1.00......0.1250 = 0.85 beam 1.00 < [b] < 2.00......0.2500 = 1.70 beam


Velocity coverage: 348.00< 1 <353.50.........-244 to 88 km/s

353.50< 1 <355.50.........-179 to 153 km/s

355.50< 1 < 5.00.......-319 to 319 km/s

5.00< 1 < 13.00........-88 to 244 km/s







29


TABLE 2.3


CO-FREE REFERENCE POSITIONS


1 b 1 b

(0) (0) (0) (0)


347.00 348.00 348.00 352.00 353.00 353.00

354.00 355.00 355.00 356.04 357.00 357.00 359.80

0.00 0.00


3.00

-5.00

4.00

-5.00

-5.00

4.00

-5.00

5.00 5.00

-5.23

-5.00

3.00

-5.75

-4.00

-3.00


0.70 1.00 1.90 3.80

4.75 5.00 7.00 8.00 8.50 9.00

10.00 11.00

12.00

14.00 15.00


-3.30

-3.90

-5.35

-4.80

2.50

-4.00

5.00

-5.00

2.75

5.00 5.00

-4.00

5.00 2.50

-5.00







30


TABLE 2.4


MOLECULAR MASS IN THE GALACTIC CENTER


Reference Observations Area Mass

lxb pc 108 M0




Audouze 1979 12CO,13CO,C180 600x600 1.3


Heiligman 1982 13CO 340x140 1.2


Sanders et al. 1984 12CO 3000x320 3-5


Bania 1986 12CO, 13CO 550x90 3.3


This work 1987 12CO 490x190 2.7























RESEIVER












TACHOMETER

MOTOR
BRAKE
. ENCODER


Fig. II-1 Diagram of the Columbia Southern
millimeter-wave Telescope.


31


E EN RS ENCR



MOTOR BRAKE TCmOMETER








32


10



6








Cr






0


-300 -200 -100 0 t00
LSR VELOCITY [KM/S]


200


300


Fig. I1-2 12CO spectrum in the direction of the Galactic
center (l,b=00,00) formed by combining two
overlaping spectra, each 333 km s- wide.


I , I I ,I I I










33


0.0


8
2.000 1.750 1.500

1.250 1.000 0.875 0.750 0.825 0.500 0.375

0.250

0.125

0.0

-0.125

-0.250

-0.375

-0.500

-0.625

-0.750

-0.875

-1.000

-1.250

-1.500

-1.750

-2.000


LSR RADIRL VELOCITY (KM S'








Fig. 11-3 Representative survey spectra; the baselines are

spaced by 6 K.


L - 0.125


--l.A


A










--A












.A


-AI




A.


I j i I i a a i I i i i i I i i 1 t i i i i i i i i i i i i i


-






















Fig. 11-4


Distribution of the velocity integrated CO antenna temperature in the plane of the sky. ThT same data, integrated between V=-320 and 320 km s are displayed at different contour levels. The contour intervals and first intervals have been set at 13, 32.5 and 65 km s in the top, middle and bottom maps, respectively. This corresponds approximately to 4, 10, and 20 standard deviations, respectively.



















00



-2*


LiJ





C
-J



< -20 .-J
~4 CD


BEAM



0I 5




-20


00

GALACTIC LONGITUDE


-50


-100


(U,


- I * 1 1






, c. c


0~ -xi

-V


10 0


50








36


-n


6000 5000





3000


2000 1000


01


10


TACTIC LONGITUDE Ides


-to


Fig. 11-5 Longitude profile (double in egral) integrated
between V=-320 and 320 km s- , and
between b=-10 and 10.


I -''N

















- . a I . . , , i . . , , I . , , , I . , , , I , ,








37


7000 6M0j


5000k-


GALACTIC LONGITUDE ( de i,


Fig. 11-6


Longitude profile integrated between V=-320 and 320 km s~ , and from b=-10 to 00.


2000 2000


0


10


-to


. . . . 1 . . . , i . . . . 1 . , , , i . .









38


5000 I-


2000' 1000


0


10


0 -s
APLACTIC LONGITUK ( dee. I


-10


Fig. 11-7 Longitude profile integrated between
V=-2O and 320 km s , and from
b=0 to 10.


-o
-u


p I I~


I -






















13.0








8.0 3.0








-2.0








-7.0








-12.0


Fig. 11-8


. %) 1\ I Ioj .W r)N r "I-- - ---orN 0000--zDONvacQO 0o oGM OOOOOO - * OOOOOOOOOOO O0 000 * OD P F *F
* 00 0 0 0 CD +D 0 C C OD OO
O O 000 C 00O0



L-V diagram integrated between latitudes
b=-2.50 and 2.50. Contours are set at
2 K from 2 K.


39


I


p

i El t il I||11i Ii I I I


]14 104 I II i 1 1 1 1 1 1 1 1 1




















I I I I I


13.0 8.0 3.0






-2.0






-7.0






-12.0


I II|


















-


3




5


2 1\2











11
''' 11''


I I I I I I I I I I I I I I I O N s O G --" -- -- -- --0
wr.J9NtNNJ--------------- e:N ooooo Qo a
OI~uJO C QONIVOOO* * * 00000000000
00000000000 -* * . 0 00......................
00000000000



Fig. 11-9 Squematic 1-V diagram, the numbered structures
are identified in the text.


40


' ' ' '


13 14
15








9



2
7 t+
6 1



4


II


II
















CHAPTER III
WIDE-LINE CLOUDS NEAR THE GALACTIC CENTER


The most remarkable objects mapped in our survey are the wide-line features located at (l,b) = (3.20, 0.30) and (5.30, -0.30). Their respective CO luminosities exceed 8% and 3% of the CO emission of the inner 500 pc of the Galaxy, integrated over all observed latitudes and velocities (see Fig. 11-4). Thus, the wide-line objects are the most intense localized CO sources in the Galaxy beyond the nuclear region.

The velocity spans of these outstanding features are the greatest observed for molecular gas in the Galaxy outside the nuclear disk. The velocity widths (FWHM) reach 140 km s-1 for the object at 1= 3.20 and 100 km s1 for the object at 1= 5.30 Their CO emission is remarkably smooth along the large velocity spans, even though some structure is also apparent (Fig. III-1).

The unusual characteristics of these objects are particularly notable in 1-V diagrams (Fig. 111-2),

integrated over 10 in latitude below (a) and above the Galactic plane (b). The wideline objects appear as intense horizontal emission streaks at 1= 3.20 and 5.30. Contrary to their large velocity extents, their longitude widths are


41







42

small, confined to about 0.50 (4 beamwidths). This apparent

"longitude-crowding" morphology that the wide-line objects present in the 1-V plane contrasts sharply with the 1-V signature of known arm-like structures, which span up to tens of degrees in Galactic longitude and usually have velocity widths well under 40 km s-1.

Although previous observations of the central region of

the Galaxy in OH (McGee et al. 1970) and CO (Bania 1977) have offered some evidence of peculiar "clumps" or "features" with surprisingly wide lines, large portions of their emission were missed because of undersampling. Owing

to the improved sampling density and latitude coverage of our survey we have been able to determine, for the first time, the total extent and CO luminosity of these objects, thus revealing their real importance (Bitran et al. 1985).

Other, less prominent wide-line features can be seen at (1, b, V) = (-5.30, 0.40, 84 km s~1); (-4.40, 0.60, 72 km s~ ); and at (-3.80, 0.90, -83 km s-l) (Fig. 111-2). It is possible that the objects at 1= 3.20 and 5.30 are the largest members of a hierarchy of wide-line clouds in the

inner Galaxy, some of which may even be concealed by the complex emission of the nuclear region, for example at (1, b)= (1.20, 0.30) (Fig. 11-4).







43


Kinematic Features or Molecular Clouds?

The interpretation of these unique objects poses several difficult questions, and startling possible consequences. The first question that arises when pondering the nature of the wide-line features is whether they are: a) kinematic features-- superpositions of many clouds along the line of sight, like the tangent point of a spiral arm--or b) real objects--localized condensations like molecular clouds

found elsewhere in the Galaxy, whose structure along the line of sight is similar to that observed in the plane of the sky.

The first interpretation has been adopted by Bania et al. (1986) and by Stark and Bania (1986), who have observed

two of our wide-line clouds ((1,b) = (3.20, 0.30) and (-5.30, 0.40)) at high resolution in 12CO, 13CO and CS. Bania et al. (1986) resolve the cloud at (l,b) = (-5.30, 0.40) into 3 molecular complexes with smaller velocity widths. They argue that the 3 components are not bound to

each other and attribute their clustering in space and velocity to gravitational forces working on a galactic scale. As for the wide-line cloud located at (l,b) = (3.20,

0.30), observations by Stark and Bania (1986) resolve the cloud into 16 components. They argue that the 16 components are aligned along the line of sight forming a spiral arm, or dust lane in the inner Galaxy.







44

Although possible, we consider kinematic explanations for the wide-line clouds extremely unlikely. A kinematic interpretation is most unlikely for a collection of wideline clouds, since it would imply an extremely privileged location for the observer. It is also an unlikely explanation for an individual wide-line object because, in addition of requiring alignment of several clouds along the

line of sight, the straight horizontal tracks of the wideline clouds in the 1-V diagrams (Fig. 111-2) do not resemble

the 1-V shapes of recognized kinematic features, such as, for example, the tangent point of the Carina arm. While the

tangent point of the Carina arm is symmetrical about the plane in b, it is highly asymmetrical in 1, with a well defined boundary in one direction (at the arm's edge) and diffuse at the opposite direction (Cohen et al. 1985). In contrast, the wide-line objects have no symmetry about the Galactic plane but are symmetrical in 1 with well defined boundaries in both directions.

Almost conclusive evidence against a kinematic

interpretation, moreover, is provided by the nearly complete

abscence of a 21 cm counterpart to either wide-line cloud (see e.g. Sinha 1979), since kinematic structures such as spiral arms, tangent points, or dust lanes are characteristically conspicuous in H I emission.

Based on the above discussion, we conclude that the wide-line features are most likely to be high-density







45

localized clouds similar to molecular clouds found elsewhere

in the Galaxy, albeit peculiar because of their large internal motions. This conclusion, we emphasize, does not preclude internal structure, which is generally characteristic of large molecular clouds and cloud complexes in the Galaxy (see e.g. Dame 1984).

Because the wide-line clouds are closely confined to the Galactic plane, cluster about the Galactic center, and

have high radial velocities, they are almost certainly a population associated with the Galactic center. Therefore,

they must be located not much further from the Galactic nucleus than their projected galactocentric distances, 0.4 to 0.8 kpc.

The radii, masses, and energies of internal motions for the wide-line clouds have been calculated on the assumption

that they are molecular clouds located at the Galactic center (Ro=8.5 kpc). These values, along with other observational parameters, are listed in Table III-1.





The Masses

Owing to their high intensity and fairly sharp edges,

the wide-line clouds are defined with little ambiguity against the CO background emission of the inner Galaxy. After subtraction of background emission in WCO maps integrated over the relevant velocity ranges, effective







46


radii (r) are defined so that nr2 is the area within the WCO contour at 1/e of peak intensity. Molecular masses, MCO, are derived from the WCO summed over the effective area of

each cloud using the conversion factor (Bloemen et al. 1986),


(N H2 = 2.8 x 1020 mol cm-2 K km- s (3.1)
WCO


and a mean molecular weight of 2.76 x 10-24 gm (Allen 1973). A detailed description of the calculations of the masses is

found in Appendix A, and their values are listed in Table III-1.

It is important to note, however, that using the above conversion factor, which was derived from gamma-ray and CO observations of the outer Galaxy, may lead to an

overestimation of the masses of clouds near the Galactic center by as much as a factor of 10 (see "Molecular Masses near the Galactic Center" in Appendix A). Therefore, Table III-1 also lists the lower limit for the masses of the wideline clouds, MCO 0.

Because of their extraordinary characteristics, the two

leading wide-line clouds at 1=3.20 and 1=5.30 were also fully mapped in 13CO. We found that the clouds' line intensity ratio 13CO/12CO was approximately 1/10, or about half of the typical Galactic value. Integrated 12CO and 13CO spectra, summed over the angular extent of the clouds







47


are displayed in Fig. 111-3 for the cloud at (l,b)=(3.20,0.30), and in Fig. 111-4 for the cloud at (1,b)=(5.30,-0.30). Using the spatially coincident 13CO and 12CO observations and assuming local thermodynamic equilibrium (LTE), masses for these clouds were calculated following the LTE method outlined in Appendix A. The values for the LTE masses, MLTE = 5 x 106 MO for the 1=3.20 cloud and MLTE = 1 x 106 MO for the 1=5.30 cloud, are within the range of the masses calculated from 12CO luminosities in Table III-1, but closer to the lower limit (MCo/lO).

The virial masses of the wide-line clouds, calculated

on the assumption of spherical symmetry with uniform mass distribution and using the observed velocity dispersions (see virial method in Appendix A), are an order of magnitude

larger then the masses derived from the CO luminosities (Table III-1). If the standard value for the conversion factor N(H2)/WCO indeed overestimated the CO masses, the discrepancies are even greater. The 3.20 and 5.30 clouds seem to be out of virial equilibrium, apparently expanding

almost freely on a time scale t (calculated as the radius divided by the rms velocity width) of 0.7 x 106 years. Calculations for the smaller wide-line objects give similar results with time scales of approximately 1 x 106 yr (Table III-1).







48


Origin of the Energy

The kinetic energies of internal motions of the two leading wide-line clouds, calculated assuming that they are

spherically symmetric, homogeneous, and expanding, range between 1052 and 1054 ergs (Table III-1). A source of energy capable of producing about 1054 ergs in a small region (r-50pc) is required to drive the largest wide-line cloud. Such high energies are unlikely to be produced by supernovae, stellar winds, cloud collisions or tidal forces,

which are the main energy sources for the interstellar medium (ISM).

Let us consider first stellar energy sources. The

total energy output of a supernova outburst is estimated to be -1053 ergs. The kinetic energy of the ejecta is about 1051 ergs (Chevalier 1977), following the explosion this energy is transferred to the interstellar medium. Stellar winds from 0 stars are also an important source of energy for the ISM; the stellar wind of a massive 0 star (M > 20 MO) can impart, over the lifetime of the star, a mechanical energy to the ISM that is comparable to that of 'a supernova outburst (McCray and Snow 1979). Assuming then, that both a supernova and an 0 star each deposit at the most - 1051 ergs of kinetic energy in the parental molecular cloud, at

least 102 to 103 such events are required to drive the largest wide-line cloud.







49


Starbursts producing large number of 0 stars and

supernovae have been shown to account fairly well for the characteristics of active energetic sources such as the nuclei of M82 and NGC253 (Rieke et al. 1980). In the case of the wide-line clouds, however, the copious ionization and far-infrared emission that such massive star formation would produce, are almost completely absent. The main wide-line clouds are inconspicuous in the IRAS and other far-IR surveys of the Galactic center (e.g. Campbell et al. 1985), practically invisible in radiocontinuum emission (Altenhoff et al. 1978), and only a few weak 5 GHz sources lie in their vicinity (Downes et al. 1980). The deficiency in ionization

is demonstrated particularly well by a comparison of the weak 5 GHz continuum sources found in the vicinity of the 3.20 cloud (Altenhoff et al. 1978) with W49, the giant H II

region at 14 kpc resulting from approximately 10 0 stars. The only continuum sources within 10 of the 3.20 cloud are

at least one order of magintude less intense than W49 and are, therefore, presumably the result of only a few 0 stars or supernovae. On the hypothesis of starburst, therefore, the ionization in the vicinity of the 3.20 and 5.30 clouds would be deficient by 2 to 3 orders of magnitude, and the far-IR emission by at least an order of magnitude. We conclude that, because of the lack of ionization and far-IR radiation, 0 stars and supernovae cannot account for the kinetic energy of the leading wide-line objects.







50


Out-of-plane neutral hydrogen observations provide

evidence of a population of high-velocity clouds (HVC), distributed through all galactic longitudes outside the disk, with velocities in excess of 80 km s1 (see review by

Hulsbosch 1975). In a Galactic fountain model, Bregman (1980) proposed that the HVC condense from a Galactic corona formed by supernova-heated gas from the disk and

subsequently fall back onto the disk. A massive cloud falling from the Galactic halo onto the disk, where it collides with molecular clouds, can transfer large amounts of kinetic energy to a small region in the plane. However, the energy transferred to clouds in the disk by a falling

HVC would be several orders of magnitude lower than that required to drive the leading wide-line clouds, since the masses and velocities of the HVC are too small. Observed velocities of the HVC rarely exceed 200 km s-1, and although their masses and distances are poorly known, the most widely accepted view is that the HVC are only a few kpc out of the

plane of the Galaxy (Hulsbosch 1975); therefore, their masses probably range between 102 and 104 M0. Moreover, the large disturbances that collisions of the needed magnitude would cause in the disk are not apparent in our CO data.

The possibility that the wide-line profiles were caused

by Galactic tidal forces disrupting the clouds was also investigated. A necessary condition for a cloud to be stable against tidal disruption is that its self-gravitation







51

should be greater than the Galactic tidal forces acting on the cloud. For a spherical cloud at galactocentric distance

ro, the minimum average density for which the cloud is gravitationally bound is given by (Mihalas and Routly 1968):


3W
PC 2nG W0 - Id/drr] (3.2)


where G is the gravitational constant, Wo = 0/re is the average angular velocity, and dO/dr is obtained from a model

for the Galactic rotation curve 0(r) (Burton and Gordon 1978). Expressing the mass in solar masses, and the volume in cubic parsecs, the critical density can be written as


M
C = 1.11 x 104 W0 ( - Ido/drr 3 'pc (3.3)



The critical densities for the wide-line clouds,

calculated with the above formula, are listed together with the observed densities in Table 111-2. Comparison of the corresponding density values indicate that, if the standard

conversion factor between CO luminosity and H2 column density (Appendix A) is valid in the neighborhood of the Galactic center, the wide-line clouds are stable against tidal disruption. On the other hand, if the CO masses of the clouds near the center have indeed been overestimated by

a factor of ten as proposed by Oort (1977), the wide-line clouds would be unstable against tidal disruption. In case







52

that this holds true, tidal disruption would be expected to result in elongation and/or break-up of the clouds radially

towards the Galactic center, i.e., the tidally disrupted clouds would appear elongated in Galactic longitude but not in latitude. However, exactly the opposite is observed in

the wide-line clouds; they are more extended in Galactic latitude than in longitude. For these reasons it is unlikely that Galactic tidal disruption plays an important role in the production of the wide-line profiles.



Are the Wide-Line Clouds Powered by the Galactic Nucleus?

The inability of conventional energy sources to produce enough energy to power the largest wide-line clouds prompts

the search for more exotic possibilities. The central position of the wide-line clouds and the magnitude of the energies involved, point to the Galactic nucleus as a likely energy source.

Considering the direct observational evidence now

accumulating for energetic phenomena at the Galactic nucleus (Lo 1985) and for resemblance to a Seyfert nucleus (Kassim et al. 1985, 1986), the general explanation of the wide-line

clouds that we find most plausible is that they are the large molecular clouds that happened to be in the vicinity of the Galactic nucleus at the time of a Seyfert outburst about one million years ago.







53

The bolometric luminosity of a Seyfert nucleus ranges from 1043 to 1046 erg s-1 (Wilson 1982); a typical value being 3.8 x 1045 erg s~1 (~1012 L0) according to Osterbrock (1984). The time scale of a Seyfert event is estimated from the characteristic velocities (103 to 104 km s-1) and sizes (-1 kpc) of the disturbed central regions to be between 105 and 106 years. A moderately bright Seyfert event with a luminosity of ~10 11 Lo (7.6 x 1044 ergs s1 ), and lasting an average time of 5 x 105 yrs, will deposit in the surrounding ISM a total energy of 1.2 x 1058 ergs. The largest wide-line cloud subtends an angle of about 10-2 sr as seen from the Galactic center, and in such Seyfert outburst at the Galactic nucleus would be illuminated by ~1056 ergs of radiant energy. A fraction of 1% of this energy transformed to kinetic energy would account for the observed internal motions.

In the Seyfert outburst scenario, the large internal motions of the wide-line clouds would be the result of rapid ionization and heating "in situ" by the UV and X-rays of the flat Seyfert spectrum (and possibly by relativistic

particles) followed by isothermal free expansion of the resultant hot (-2x104 K), high-pressure fireball through the sonic speed (-10 km s- ). During the outburst, a distant observer might have classified the cloud as a Seyfert forbidden-line filament. In fact the velocity







54

widths, sizes, and densities of the forbidden-line filaments observed in Seyfert Galaxies (Osterbrock 1984) are comparable to those of the wide-line clouds. On termination

of the outburst, without a sustained source of ionizing photons, the electrons and ions will recombine and become neutral in a short time. Lo (1986) quotes a recombination time of less than 10 years for Sgr A West. The

recombination time is inversely proportional to the ionized gas density and, therefore, it is likely to be about a few hundred years for the wide-line clouds. Molecular formation would follow on a time scale difficult to specify (since it would depend on the fate of the dust grains), but presumably well within the last million years.

Once set in motion, the dense gas of a large molecular cloud is not readily stopped, as Oort (1977) has emphasized,

and the wide-line objects would be among the most durable remnants of a Seyfert outburst in a spiral Galaxy such as ours. The 3 kpc and the 135 km s~1 classic expanding arms

of the inner Galaxy could be descendants of a previous generation of wide-line objects (the masses and the kinetic

energies are about right, if some swept-up material is added), and the difference in age between these arms and the

wide-line objects would be a rough measure of the time between successive outburst: about 10 Myr.

Although the proposed mechanism can account for the

high energies needed to drive the wide-line clouds and could







55

also explain some of the expanding features farther from the

nucleus, it has several weak points: It is, of course, debatable that our Galaxy has had Seyfert-like behavior, with an outburst taking place as recently as a million years

ago; the mechanisms for the absorption of energy by the clouds and the efficiency of transformation to kinetic energy are unclear; the velocity widths of the wide-line clouds are a few times smaller than those observed in Seyfert filaments and do not surpass the velocities allowed by the rotation curve of the Galaxy. Also, as a consequence of the outburst, the population of wide-line clouds near the Galactic center would be expected to be larger. This may indeed be the case, however, since wide-line clouds located

in the central 30 of the Galaxy are hard to recognize because they would blend with the complex emission of the region. In fact, in 1-V diagrams (Fig. 11-8, 111-2), the central region appears studded by several horizontal streaks which could be unrecognized wide-line clouds.

In summary, the wide-line features are among the most

puzzling phenomena observed in the central region of the Galaxy. Their characteristics suggest that they are

localized high-density molecular clouds situated near the Galactic center. Conventional mechanisms cannot provide the

energies required to explain their large internal motions. We propose that the Galactic nucleus is the energy source for the wide-line clouds, which implies that our Galaxy







56

underwent a Seyfert-like outburst about one million years ago.







57


TABLE 3.1


WIDE-LINE MOLECULAR CLOUDS


Object (a) (b) (c) (d) (e) (f) (g) (h)

10, b0




5.3,-0.3 95 43 786 31 0.7-6.4 64 35-350 0.7

3.2, 0.3 104 60 477 43 2.3-23 175 250-2500 0.7

-3.8,0.9 -83 19 579 19 0.2-1.7 8 2-18 1.0

-4.4,0.6 72 22 658 20 0.2-1.7 11 3-25 0.9

-5.3,0.4 84 20 787 25 0.5-5.2 12 6-62 1.2


(a):

(b):

(c):

(d):

(e):

(f):

(g):

(h):


cloud mean velocity [km/s] rms velocity width, AVrms [km/si projected distance from Galactic nucleus [pci effective radius [pci CO mass range [106 M0] virial mass [106 M0] internal kinetic energy range [1051 ergs] characteristic time (r/AVrms) [106 yr]








58


TABLE 3.2


DENSITIES OF THE WIDE-LINE CLOUDS


Object (a) (b) (c)

l(0), b(0)




5.3, -0.3 71 7.1 13.6

3.2, 0.3 96 9.6 33.0

-3.8, 0.9 80 8.0 25.2

-4.4, 0.6 69 6.9 20.1

-5.3, 0.4 108 10.8 13.5


(a): measured cloud density

(b): lower density limit

(c): critical density


[M0/pc3 [M0/pc3 [M0/pc3








59


J 3.250 Ii I \ I

1.000




0.75* 0.500 0.25*




0.00*




- 0.250




-0.500




-0.75* 5K




- 1.00*

SI I I
-200 0 200

VLSR(km s-)


Fig. III-1 Spectra taken across the 1=3.20 cloud.





















Fig. 111-2 Longitude-velocity diagrams integrated over latitude
from b=-l to -0.125 (left) and from b=0.1250 to 1.00
(right). Contours are set at 0.125 K-degree intervals
with the first contour at 3 standard deviations. The
arrows point to the 2 largest wide-line clouds.

















b<00


























Resolution


I -


-


10' 56




0
-j 0






3556 3500


-200 -100 0 100 200 300 -300 -200 -100


r b >0*


I- i


0 100 200 300


LSR RADIAL VELOCITY (km s)


a'N


-I


-300


-0


























- Resolution


-



















I I--









62


100 Uj S80 E 60



0


20



0


-300 -200 -100 0 100
LSR VELOCITY [KM/SJ


200


Fig. 111-3 12CO and 13CO spectra of the wide-line cloud
at (l,b)=(3.20,0.30), summed over the cloud's
angular size.


- I I I I I -









63


-300


-200


100 0
LSR VELOCITY


100
[KM/S)


300


Fig. 111-4 12CO and 13CO0 spectra of the wide-line cloud
at (l,b)=(5.30,-0.30), summed over the cloud's
angular size.


-- 50
-so


40 30


20


S10


0


, * , , , I i, II , , , I , , ,l , i I I
















CHAPTER IV
A CATALOGUE OF MOLECULAR CLOUDS IN THE INNER GALAXY


A distinctive characteristic of the molecular gas in the Galaxy, in contrast with the atomic gas, is its clumpy spatial distribution. While the atomic gas is distributed smoothly, the molecular gas is found concentrated in clouds and cloud complexes. The larger complexes, self-gravitating aggregates with masses of the order of 105 Mo and radii of approximately 50 pc, are known as "giant molecular clouds" (GMC). They are recognized as the phase of the Inter Stellar Medium (ISM) where most star formation takes place and are useful tracers of spiral structure (Dame 1984).

Because of its clumpy nature the distribution of molecular gas in the Galaxy is best quantified by an inventory or catalogue of molecular clouds. A catalogue that lists the location and physical characteristics of molecular clouds in the inner Galaxy is useful to study spiral structure and star formation in this region, and for comparison with molecular clouds in the outer Galaxy. In this chapter we present a catalogue of molecular clouds in the inner Galaxy. The catalogue lists the mean coordinates,

radial velocities, radius, velocity width and apparent CO luminosity of the molecular clouds in the region (see Tables


64







65

IV-1 to 8). This is the only catalogue of molecular clouds

in the inner Galaxy, available to date, with a complete latitude coverage.

The catalogue, which includes the observed highvelocity clouds, is organized by membership of the clouds in recognized H I structures such as the 3 kpc arm, the 135 km s51 arm, etc. (see Chapter II). The membership of the clouds in these structures was decided by spatial and velocity coincidence of the clouds with the locus of the H I structures as they appear in a survey by Cohen (1975). In

several instances, superpositions of features make the identification of clouds and their membership in H I structures uncertain. Therefore, it is likely that the classification of some of the clouds may have to be revised when higher-resolution observations become available. The

clouds at the Galactic center are difficult to identify against the intense, wide-velocity background, therefore they have not been included in the catalogue. The wide-line clouds are catalogued in chapter III.

The procedure for determining the boundaries of the clouds was done as follows. The clouds were identified in "clipped" latitude-integrated 1-V diagrams. These diagrams

are produced by setting to zero the contributions of channels with antenna temperature lower than a certain threshold value. The molecular clouds are seen distinctly

in this type of diagram, when the threshold has been







66


properly set to eliminate the difusse background emission. Through trial and error, a threshold temperature of 2 K was

found generally adequate to produce good contrast in the diagrams, without eliminating identifiable clouds. An

integrated spectrum (T vs. V) was then produced for each cloud by adding the spectra within the spatial boundaries of the cloud. The boundaries of the clouds were found by inspection of the 1-V, 1-b, and b-V diagrams displayed in appendices B, C, and D. The limits of the clouds were revised by examining temperature profiles, and the procedure iterated until all the emission from the cloud was included in the integrated spectrum.

The location and velocities of the clouds in the catalogue are given as the intensity-weighted mean coordinates. We have chosen these coordinates rather than the more commonly used, and easier to find, peak coordinates because they are more physically meaningful. The intensityweighted coordinates are closer to the center-of-mass coordinates of the cloud, since the CO intensity is proportional to the molecular mass (see Appendix A)

The intensity-weighted mean longitude was calculated as,


T. (l,b,V) 1
< 1 > - l1,b,V
Y T. (l,b,V)
l,b,V


similarly the mean latitude is,







67


1 T~ (1,b,V) b(42
< b > = l,b,Ji (4.2)



and the mean velocity,


T (l,b,V) V
< V > = TibVV (4.3)
ZT. (1,b,V)
l,b,V



The listed radius corresponds to the equivalent radius

of a circle having an area equal to the cloud's observed calclatd asr =1/2
area, and is calculated as r = (A/n) . The velocity

widths of the clouds, found in Tables IV 1-8, correspond to the FWHM calculated from Gaussian analysis of the integrated spectra. The integrated spectra were also used to calculate the apparent CO luminosity, ICO. This quantity was

calculated according to the expression: ICO = Tsum x Dl x Db x DV, where the area under the integrated spectrum is Tsum,

the spatial picture elements are Dl=Db=0.1250, and the velocity resolution element is DV=1.3 km s-1. The apparent

CO luminosities are followed by their corresponding rms uncertainties.

Because of the large noncircular motions observed in the central region of the Galaxy, kinematic distances to clouds in this area are not reliable; therefore they are not listed in the catalogue. For some clouds, however, it has been possible to find associated H II regions by spatial and







68


velocity identification in catalogues of H II regions (e.g. Georgelin and Georgelin (1976), Lockman (1979), and Downes et al.(1980)). When an estimate for the distance of the associated H II region was available, the molecular mass of

the clouds was calculated. The mass of a cloud can be calculated directly, in solar masses, from the CO luminosity ICO, and the distance d in kpc, with the formula


M = 1.9 x 102 CO d (4.4)


For details on the mass calculations see Appendix A.




Notes on Selected Clouds.

We have examined catalogues of H II regions looking for

counterparts to the 'molecular clouds listed in the catalogue. Associations of molecular clouds and H II regions provide evidence of star formation and of real density enhancements as opposed to features produced by velocity crowding. The identification of molecular clouds

with H II regions is far from exhaustive because of incompleteness of the available H II catalogues in spatial and velocity coverage.

Cloud # 15, 3 kpc arm. This cloud located at (1, b, V) = (5.350, -0.030, -20 km s-l) may be associated to the H II region at (1, b, V) = (5.4790, -0.2410, -26 km s~1) (Downes et al. 1980). This association would provide substantial evidence of star formation in the 3 kpc arm, contradicting







69

earlier observations (Lockman 1980) in which no tracers of star formation were found. If this cloud is located at a distance of 5.5 kpc, its molecular mass would be 4.7 x 105 M0.

Cloud # 3, III b. Cloud # 3 at (1, b, V) = (7.560,

0.050, 122 km s-1) is probably associated to the H II region at (1, b, V) = (7.4720, 0.060, 121 km s-1). Downes et al. (1980), argue that this source is located at 27 5 kpc. This distance implies a molecular mass of 3+1 x 106 Mo.

Cloud # 8, III b. The location of this cloud, (1, b, V) = (12.170, -0.020, 118 km s-1) indicates an association with a source at (1, b, V) = (12.2840, -0.1160, 119 km s- ). The distance in Downes et al. catalogue, 16 1 kpc, implies a mass of 1.2 x 106 M0.







70


TABLE 4.1


MOLECULAR CLOUDS IN THE 3 KPC ARM


<1> V r ICO

(0) (0) (km/s) (km/s) (') (a) (a)


1 -10.24 2 -9.37 3 -8.47 4 -7.79 5 -7.46 6 -6.66 7 -5.87 8 -4.49 9 -3.29 10 3.52 11 4.28 12 4.76 13 5.35 14 9.05 15 10.56


0.13

0.43

-0.29

-0.14

-0.11

0.08 0.13 0.19 0.29

-0.49

0.45 0.24

-0.03

-0.13

-0.37


-107 10 15 4.18


-105

- 95

- 93

- 84

- 87

- 85

- 76

- 70

- 28

- 28

- 28

- 20

- 15

- 2


13

8

6 10 11

9 13

7 11

9

8

8

8 10


15 25

22 25 31

21 16 13 23

12 12 25 19 17


6.33 7.89

7.43

13.25

21.49 5.24 7.49

4.13 9.43 3.21 3.29 6.93

7.49 8.17


(a) = [K sq.


deg. km/sI


* = seem to form part of a larger complex


0.03

0.04

0.04 * 0.04 * 0.05 * 0.07 *

0.04 0.03

0.02 0.04 0.07

0.02 0.04 0.04 0.03







71


TABLE 4.2


MOLECULAR CLOUDS IN THE CONNECTING ARM


<1> V r ICO C3

(0) (0) (km/s) (km/s) (') (a) (a)




1 3.32 -0.52 260 27 15 8.46 0.05

2 3.98 -0.67 237 24 23 10.49 0.08

3 4.83 -0.87 206 29 27 22.3 0.1

4 5.45 -0.74 194 24 17 6.52 0.05

5 6.59 -0.68 158 22 34 23.7 0.1

6 7.77 -0.64 142 27 34 43.8 0.1

7 9.35 -0.91 120 16 40 11.3 0.1


(a) = [K sq. deg. km/si







72


TABLE 4.3


MOLECULAR CLOUDS IN FEATURE III b


<1> V r Ico G(0) (0) (km/s) (km/s) (') (a) (a)




1 4.33 0.01 89 10 12 1.99 0.03

2 6.76 0.07 118 12 12 1.06 0.03

3 7.56 0.05 122 10 13 2.30 0.03

4 8.32 0.00 126 15 13 3.83 0.04

5 9.20 -0.02 126 10 17 4.21 0.04

6 11.54 0.01 120 10 10 1.23 0.02

7 12.17 -0.02 118 17 13 2.46 0.03

8 12.80 0.01 101 8 10 0.78 0.02


(a) = [K sq. deg. km/s]







73


TABLE 4.4


MOLECULAR CLOUDS IN THE 135 KM/S ARM


# <1> V r ICO

(0) (0) (km/s) (km/s) (') (a) (a)




1 -0.63 0.04 132 56 25 58.0 0.1

2 2.08 -0.15 133 41 24 36.1 0.1

3 2.45 0.29 159 15 15 3.67 0.04

4 5.34 0.26 182 7 22 2.60 0.04

5 5.89 0.33 190 10 15 1.10 0.03

6 7.92 0.29 183 15 19 3.99 0.05

7 8.94 0.36 181 17 22 4.06 0.06


(a) = [K sq. deg. km/s]







74


TABLE 4.5


MOLECULAR CLOUDS IN FEATURE IX


# <1> V r ICO G

(0) (0) (km/s) (km/s) (' (a) (a)




1 5.92 0.01 137 12 13 3.15 0.04

2 6.60 0.03 141 5 9 0.43 0.01

3 6.70 -0.03 152 7 12 0.90 0.02

4 7.69 -0.06 153 7 12 1.21 0.02

5 10.65 -0.01 153 9 18 2.96 0.04

6 12.19 -0.08 147 10 13 0.96 0.03

7 12.73 -0.04 145 12 9 0.74 0.02


(a) = [K sq. deg. km/s]







75


TABLE 4.6


MOLECULAR CLOUDS IN THE NORMA ARM


# <1> V r ICo C(0) (0) (km/s) (km/s) (') (a) (a)




1 -11.30 0.165 -68 17 13 7.98 0.04 * 2 -10.82 0.130 -66 15 13 7.41 0.04 * 3 -9.95 0.082 -64 17 18 11.85 0.05 *

4 -8.49 0.252 -45 7 17 4.82 0.02 5 -8.00 0.-185 -53 12 15 5.40 0.03 6 -6.69 -0.040 -49 20 22 13.41 0.05 7 -4.36 0.690 -34 7 17 4.08 0.03 8 -1.99 -0.400 -24 7 17 5.47 0.03


* seem to form part of a larger complex

(a) = [K sq. deg. km/s]







76


TABLE 4.7


MOLECULAR CLOUDS IN THE SCUTUM-CRUX ARM


# <1> V r ICO 0

(0) (0) (km/s) (km/s) (') (a) (a)




1 -10.11 0.157 -39 12 21 10.54 0.04

2 -9.19 0.414 -36 15 13 4.88 0.03

3 -5.24 0.640 -21 7 17 10.94 0.04


(a) = [K sq. deg. km/s]







77


TABLE 4.8


MOLECULAR CLOUDS IN THE 70 KM/S ARM


# <1> V r ICo 0(0) (0) (km/s) (km/s) (') (a) (a)




1 -11.58 0.02 4 12 10 1.29 0.02 2 -10.79 -0.01 15 10 10 1.55 0.02 3 -9.98 0.03 17 10 16 3.52 0.04 4 -8.33 0.05 42 10 10 0.66 0.03 5 6.44 -0.08 81 10 10 0.98 0.02 6 7.27 -0.03 91 10 17 2.99 0.03 7 10.99 -0.04 91 17 12 1.58 0.03 8 11.96 0.02 97 10 17 2.80 0.03


(a) = [K sq. deg. km/si















CHAPTER V
THE KINEMATICS OF THE MOLECULAR GAS IN THE INNER GALAXY



Introduction

The kinematics of the atomic and molecular gas in the

inner Galaxy are complex because of noncircular motions, large-scale asymmetries in longitude and velocity, and a spatial distribution inclined with respect to the Galactic plane. The situation is further complicated by the two-fold

distance ambiguity inside the solar circle and by the pervasive effects of streaming motions of the gas. It is not surprising, therefore, that no interpretation or model proposed to date can successfully account for all the features observed.

Most efforts towards understanding the kinematics of the inner Galaxy are based on H I data and, in general, fall into three broad categories, namely: 1) ejection models, 2) gravitational-field models, and 3) phenomenological models. Their characteristics and shortcomings are discussed below.

In ejection models, the radial motions are attributed to gas in expansion expelled from the Galactic nucleus. An ejective origin of the 3 kpc and 135 km s-1 arms (Van der Kruit 1971), for example, would require the expulsion of about 108 Mo of gas from the nucleus (Oort 1977). Such


78







79

large amounts of gas cannot be replenished by normal stellar

evolution over the lifetime of the features (- 106 yr). Although this problem could be circumvented by circulation of the clouds through the center (Oort 1977), the essential

remaining objection to ejection models is that several collimated events with various energies and occurring at different specific epochs have to be invoked to explain the observed features.

Field models assume that the radial motions of the gas are the response to asymmetries in the gravitational field at the Galactic center. In order to explain the observed

radial motions and the tilt of the gas distribution using the field hypothesis, the mass distribution at the center is

required to be both non-axisymmetric and tilted to the Galactic plane. While there is no evidence that the central star bulge--the main contributor to the mass density at the

center--deviates from axial symmetry, a relatively small systematic deviation from symmetry is possible, and could explain some of the non-circular motions observed. For example, a 10% oval distortion of the bulge can account for

the radial motions of the 3 kpc arm according to Yuan (1985). However, it is unlikely that the central star bulge

has a significant tilt with respect to the plane, because the late-type stars of which it is composed have high random-velocities and are probably well-mixed. Moreover, inclined bulges are not observed in external edge-on







80

galaxies, perhaps because the torques acting on an inclined bulge (Tremaine and Weinberg 1984) would make it unstable.

Due to these constraints it is very unlikely that field models, except perhaps triaxial models (discused below), can

account for the tilt of the gas layer and several other structures such as some jet-shaped features observed in H I,

the molecular ring (Chapter II), and the wide-line clouds (Chapter III).

Owing to the complexity of the gas motions in the Galactic center region, and the lack of knowledge of the dynamics involved, purely phenomenological models -independent of dynamical considerations-- that account for different features with a unifying assumption are useful to constrain dynamical interpretations of the inner-Galaxy gas.

For example, Burton and Liszt (1978) were able to account for several observed H I features by assuming that most of the gas at the center is confined to a disk in expansion and rotation inclined 220 to the Galactic equator. Peters (1975) was able to account for the radial motions of the 3

kpc and 135 km s~1 arms by assuming that the gas at the center moves in elliptical streamlines. Liszt and Burton (1980) have also accounted for several different H I features by assuming that the gas moves along elliptical streamlines such as those produced in a bar-like potential, but without reference to a particular mass distribution.







81


The dynamical foundation lacking in phenomenological models could be obtained if one assumes that the stellar Galactic bulge is a rotating triaxial ellipsoid.

Calculations of orbits in rotating triaxial potentials (Schwarzchild 1979, Heiligman and Schwarzchild 1979) can produce closed elliptical orbits in retrograde motion and stable elliptical orbits inclined with respect to all the planes of symmetry. These orbits could explain, in principle, the presence of gas at "forbidden velocities" and the tilted distribution of the gas at the center (Lake and Norman 1982). So far, however, triaxial models have failed

to reproduce the gas kinematics of the Galactic center. Heiligman (1982) used a triaxial potential, of the form shown to be self-consistent by Schwarzchild (1979), to calculate stable orbits intended to reproduce 13CO observations in the inner 400 pc of the Galaxy. The model, constrained to the plane, reproduces the general distribution of the gas in the 1-V plane, but not the 135 km s1 arm or the molecular ring. Vietri (1986), developed realistic self-consistent models for the Galactic potential aimed at reproducing the inclined disk of Burton and Lizst (1978). He concludes that "...it is unlikely that the 220 tilt of the H I ring at r=1.5 kpc can be explained in terms

of the tilted stable retrograde orbits in a rotating triaxial bulge," Vietri (1986, pg. 62).







82


The failure of these different types of models to account for the kinematics of the gas near the Galactic center is probably a consequence of the use of oversimplified assumptions for the explanation of an intrinsically complex situation. In this state of affairs,

it seems most useful to concentrate our efforts on accurately measuring observational parameters that can provide critical constraints for the different types of models discussed above. Owing to the high sensitivity, large spatial and velocity coverage and high sampling density of the present survey, the CO data collected in this survey are particularly suited for this purpose.



The Molecular Disk

Dynamical and phenomenological models of the Galaxy should account for the distribution and kinematics of the interstellar gas. Conversely, the characteristics of the molecular gas layer can provide important constraints on such models. In this section we discuss the measurements of the inclination and thickness of the molecular gas layer in

the inner Galaxy obtained in this survey, and their relevance to models proposed for this region.







83


The Tilt of the Molecular Disk

Since the first H I surveys of the center region,

observers have noted that the emission occurs preferentialy in opposing quadrants (1>00, b<00 and 1<00, b>00), so that

the gas distribution appears inclined to the Galactic equator. The magnitude and extent of the tilt of the gaseous disk are important quantities for models of the region. In particular, they are critical for the models proposed by Burton and Lizst (1978) and Lizst and Burton (1978, 1980), since their basic assumption is that the gas in the inner Galaxy (H I and CO) is contained in structures tilted to the Galactic plane.

Although there is consensus on the direction of the tilt, its magnitude remains controversial. A tilt of about 80 was found by Kerr (1968) for the distribution of neutral hydrogen. The rotating and expanding disk proposed by Burton and Lizst (1978) is tilted by 220 with respect to b=0 and 780 to the plane of the sky. These values were found by fitting synthetic spectra to H I observations and from the inclination of high-velocity gas in moments maps. In another model, Liszt and Burton (1980) propose that the gas at the center is in pure rotation, confined to a barlike structure inclined 13.50 to the Galactic Equator, by 700 to the plane of the sky and at an orientation angle of 41.50. These models are shown to account for H I features IIIa, IV, X, and XII (see Chapter II), and the authors claim







84

that they also explain the distribution of molecular gas at the Galactic center. However, Heiligman (1982) observing 13CO within 20 of the center, found a tilt of 70 for the high-velocity molecular gas and no appreciable tilt for the

high-intensity gas at the center. Sanders, Solomon, and Scoville (1984), also found a tilt of 70 from their 12CO observations for l>-40, after subtracting the contribution of the gas in the foreground.

The disagreement on the magnitude of the tilt of the gaseous disk at the Galactic center is due, in part, to different methods used to measure the angle of inclination.

Discrepancies that remain after allowing for the different methods used probably reflect intrinsic differences in the distribution of the atomic and molecular gas at the Galactic center.

In order to investigate how the methods used to measure the inclination affect the results, we have measured the tilt of the CO emission ridge to the Galactic plane in three different ways: a) from an 1-b map integrated over all observed velocities, i.e. from the distribution of CO emission in the plane of the sky, b) from 1-b maps integrated over selected velocity ranges that exclude the contribution of local emission, and c) from the 1-b distribution of the emission at extreme velocities.

A problem encountered when attempting to measure the inclination of the molecular disk from 1-b maps is to define







85

the structure or ridge whose inclination is to be measured.

When the location and extent of the ridge is decided by eyeball estimates, considerable error occurs. In order to diminish this uncertainty, we adopted the following

procedure, first we calculated the intensity-weighted mean latitude of the CO emission every 0.250 (2 beamwidths) in longitude, and then we fitted a least-squares intensityweighted straight line to the ridge of mean latitudes. We adopted this line as a first order representation of the CO emission in the plane of the sky, and used it to define the

tilt of the molecular disk to the Galactic equator in the plane of the sky.

Figure V-1 shows the locus of the intensity-weighted mean latitudes, integrated over all observed velocities. The mean latitude ridge can be separated into 3 regions with different characteristics. In the central region (-1.50 <1<

1.80), the ridge is virtually parallel to the Galactic plane, and the mean latitude slightly negative. At negative

longitudes the ridge runs mostly above the plane reaching its maximum latitude of b=0.4190 at 1=- 3.750. At positive

longitudes, conversely, the mean latitude ridge extends mostly below the Galactic equator reaching the minimum latitude of b=-0.3870 at 1=4.750. The large peak at 1=3.250 with mean latitude b=0.1710 corresponds to the largest wideline cloud (see Chapter III).







86

The straight line between 1=50 and l=-50 in Fig. V-1 is

the intensity-weigthed least-squares fit to the latitude ridge, in the region where the tilt is observed. The line

is inclined to the Galactic equator by 2.30, an angle smaller than all the values so far reported for the tilt of the gas distribution at the center. This value, however, should be regarded as a lower limit because it was obtained

integrating the emission over all observed velocities. Since this procedure overrepresents the contribution of the

local gas, which is not likely to be inclined in a preferential way to the Galactic equator, the result is biased towards a smaller angle.

The contribution of the inner-Galaxy gas can be separated from that of local material, in velocityintegrated 1-b maps, by a suitable choice of integration limits. Figures V-2 and V-3 present the distribution of the

high-velocity molecular gas in the plane of the sky, respectively integrated from V=-319 km s-l to V=-98 km s-1 and from V=98 km s1 to V=319 km s~ These limits were

chosen because most local-gas emission occur outside these

ranges, and because these limits allow direct comparison with H I observations (Burton and Liszt 1978). Our Figures

V-2 and V-3 are the molecular counterparts of Burton and Liszt's H I moment maps (their Figs. 2 and 3).

Comparison of Figs. V-2 and V-3 with Figs 2 and 3 of Burton and Lizst (1978) reveals that the inclination of the







87

high-velocity molecular gas to the Galactic equator in the

inner Galaxy is smaller than that of the high-velocity atomic gas. Although Burton and Lizst did not measure the inclination of H I emission to the Galactic plane directly from their figures 2 and 3, they argue that the inclination

of the H I emission deduced from their H I moment maps support the tilt angle they propose for the H I disk (220 in Burton and Lizst 1978 and 13.50 in Lizst and Burton 1980).

The variation of the intensity-weighted mean latitude of CO emission with Galactic longitude is displayed in Fig. V-4 for -319 < V <-98 km s-1 and in Fig. V-5 for 98< V <319

km s~1, in the longitude ranges where significant emission takes place.

To calculate the tilt of the molecular disk from the CO

emission moment maps we have used the same technique described above, i.e. we have fitted an intensity-weighted least-squares straight line to the ridge of intensityweighted mean latitudes (Fig. V-6). The inclination of the

lines are 80 for the negative velocities and 2.20 for the positive velocities. The smaller tilt found at positive velocities is due to the contribution of the wide-line cloud

located at (l,b)= (3.20, 0.30), which produces a large fraction of the CO emission in this velocity range.

It is possible that the wide-line cloud does not form part of the inclined molecular disk at the center. If the

contribution of the wide-line cloud is subtracted, the







88


inclination of the positive-velocity CO emission to the equator is 60; similar to that found at negative velocities,

and resulting in an average inclination (positive and negative velocities) of 70. This value is in agreement with

the angles measured by Kerr (1968), Heiligman (1982), and Sanders et al. (1984). The tilt angles calculated by these authors, therefore, have not taken into account the contribution of the wide-line cloud at (l,b) = (3.20, 0.30). Neither Kerr's nor Heiligman's observations could have taken into account the contribution of the wide-line cloud. The lack of an H I counterpart to the wide-line cloud explains

Kerr's result; while Heligman's survey, extending only to 1=20, missed the wide-line cloud for over a degree. The Survey of Sanders et al. (1984) probably missed the wideline clouds due to undersampling. Their observations were

spaced by 10 in 1, about twice the longitude extent of the largest wide-line cloud.

In either case, however, the inclination of the CO

emission to the Galactic equator calculated from moment maps

is significantly smaller than the inclination of H I emission in analogous maps as reported by Burton and Liszt

(1978) and Liszt and Burton (1980), and of CO emission as reported by Liszt and Burton (1978).

The CO and H I moment maps also differ in the latitude extent of the respective emissions and their concentration to the central region. Most CO emission is confined between







89

b=-10 and b=10 , while the H I emission extends to b=40 and b=-50. This difference reflects the intrinsically smaller latitude scale-height of CO with respect to H I, and is not due to lack of latitude coverage of the CO emission. A low resolution CO survey of the center region, between b=-80 and b=80 (Nyman et ai. 1987) confirms that essentially all CO emission from the center region falls within the latitude boundaries of the present survey.

The emission from the region within a couple of degrees

of the Galactic center is clearly distinguished from the rest of the inner Galaxy emission in the CO moment maps (Figs. V-2 and V-3). In the H I maps, on the other hand, the central region is blurred and confused with emission and absorption produced by surrounding gas. The inclinaton of

the mean-latitude ridge is smaller in the central region (Fig. V-6). We have calculated the inclination of the CO emission within 10 of the center, where the largest concentration of molecular mass is found. The inclination of the intensity-weighted least-squares fit is 2.50 for the

positive velocities and 3.60 for the negative velocities. Comparison of the tilt of this central molecular disk with

the inclinations found at larger longitudes suggests that the tilt of the CO emission to the Galactic equator increases with Galactic longitude.

The velocity limits used in the moment maps still allow some contamination by emission which does not originate at







90


the center; this affects the values found for the inclination of the molecular disk. In order to avoid any contamination, and obtain an upper limit for the tilt, only the highest velocities should be considered. The latitudes at which the highest velocities occur have been plotted as a function of Galactic longitude in Fig. V-7. The latitude distribution of the terminal velocities highlights the difference between the central region and the rest of the inner Galaxy. The high-velocity gas within ~ 1.750 (~250 pc) of the center is confined to a narrow band, parallel to the Galactic equator, at b=-0.1250. Outside this central region the latitude distribution of terminal velocities forms a band, extending from 1=-60 to 1=50, inclined to the Galactic equator by 90 for 1<00 and by 130 for 1>00. Thus, the average inclination of the molecular disk, derived from the gas at terminal velocities, is approximately 110 1 20. It should be noted that the emission at higher inclination represents a small fraction of the emission (and thus of the

molecular mass) tilted at smaller angles to the Galactic equator. This value is 40 larger than the inclination determined from the CO moment maps, and within 2.50 of the inclination of 13.50 calculated by Lizst and Burton (1980). However, the inclination deduced from the extreme velocities is significantly smaller than the 220 tilt found by Burton and Lizst (1978) and Lizst and Burton (1978). The extent of the CO tilted structure in Galactic longitude, is about 110




Full Text

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CO IN THE GALACTIC CENTER: A COMPLETE SURVEY OF CARBON MONOXIDE EMISSION IN THE INNER 4 KPC OF THE GALAXY BY MAURICIO ERNESTO, B I TRAN CARRENO 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 1987

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Copyright 1987 by Mauricio Ernesto Bitran Carreno

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This dissertation is dedicated with much love to my mother Cora C. de Bitran, to my wife Gloria, to my son Jonathan, and specially to the blessed memory of my father Raul Bitran Nachary ( Z . L. ) .

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ACKNOWLEDGMENTS This work was possible thanks to the efforts and dedication of a number of persons and institutions. To all of them I express my sincere appreciation. I thank Dr. Stephen Gottesman, my advisor at the University of Florida, for his excellent guidance, continuous support, and friendship during the course of this work. Likewise I thank Dr. Patrick Thaddeus, my coadvisor, for his stimulating scientific guidance, hospitality, and unrestricted access to the computing facilities of the Goddard Institute for Space Studies. The supervisory committee members, Drs. T. D. Carr, J. Hunter, J. Ipser, A. G. Smith, H. Smith, and W. Weltner, are gratefully thanked for their useful suggestions and comments. The project of millimeter-wave observations of the Southern skies with a one-meter class telescope was conceived and planned by Dr. Patrick Thaddeus. Its successful implementation was largely due to the efforts of the project's manager Dr. Richard Cohen, with the help of Drs. Leonardo Bronfman, and David Grabelsky. J. Montani , F. Aviles, and M. Koprucu, with dedication and competence, kept the telescope running smoothly during its operation. Profs. Jorge May and Claudio Anguita, of the University of Chile, IV

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promoted and supported the involvement of Chilean astronomers in the project. I am indebted to Dr. Leonardo Bronfman for advice, as well as for his hospitality and friendship. Many thanks go to M. Koprucu and Dr. H. Alvarez for help with the observations, and to A. Smith, Dr. R. S. Cohen, Dr. L. Erickson, and Dr. T. Dame for help and advice with the data reduction. I thank the University of Chile for support during the observations at Cerro Tololo, and the faculty of the University of Florida for support and encouragement during my graduate studies. During this period we have enjoyed the friendship and support of Drew and Jill Weisenberger, Celia Gottesman, and of fellow graduate students F. Reyes, G. Fitzgibbons, J. Webb, M. J. Taylor, and W. Cooke, in particular . My deepest gratitude goes to my parents and to my wife for their permanent love and support. Dr. Gloria Rachamin, my wife, deserves special recognition for her unfailing love, support, and encouragement, as well as for her skillful editing of the manuscript. v

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TABLE OF CONTENTS PAGE ACKNOWLEDGMENTS iv ABSTRACT viii CHAPTERS I. INTRODUCTION 1 The Center of the Galaxy 1 CO Observations of the Galactic Center . . 7 Present Work 8 II. THE SURVEY 11 Instrumentation 11 Observations 14 Large-scale Characteristics of the CO Emission 17 III. WIDE-LINE CLOUDS NEAR THE GALACTIC CENTER . . 41 Kinematic Features or Molecular Clouds? . 43 The Masses 45 Origin of the Energy 4 8 Are the Wide-Line Clouds Powered by the Galactic Nucleus? 52 IV. A CATALOGUE OF MOLECULAR CLOUDS IN THE INNER GALAXY 64 V. THE KINEMATICS OF THE MOLECULAR GAS IN THE INNER GALAXY 7 8 Introduction 78 The Molecular Disk 82 The Tilt of the Molecular Disk .... 83 The Thickness of the Molecular Disk . . 92 The Rotation Curve of the Inner Galaxy . . 96 vi

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VI. SUMMARY 120 APPENDICES A. DETERMINATION OF MOLECULAR MASS IN CLOUDS . . 124 CO Masses 124 LTE Masses 126 Virial Masses 128 Molecular Masses near the Galactic Center 129 B. LONGITUDE-LATITUDE MAPS 131 C. LONGITUDE-VELOCITY DIAGRAMS 147 D. LATITUDEVELOCITY DIAGRAMS 177 REFERENCES 229 BIOGRAPHICAL SKETCH 235 Vll

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Abstract of Dissertation Presented to the Graduate School of the University of Florida in Partial Fulfillment of the Requirements fox the Degree of Doctor of Philosophy CO IN THE GALACTIC CENTER: A COMPLETE SURVEY OF CARBON MONOXIDE EMISSION IN THE INNER 4 KPC OF THE GALAXY By Mauricio Ernesto Bitran Carreno December 1987 Chairman: Stephen T. Gottesman Cochairman: Patrick Thaddeus Major Department: Astronomy The first well-sampled, large-scale survey of ^CO (J=l-»0) emission from the inner 4 kpc of the Galaxy is presented and used to study the distribution of molecular clouds and the kinematics of the molecular gas in the inner Galaxy. The survey samples a 4° wide strip along the Galactic equator from 1=-12° to 1=13°. The over 8000 spectra obtained with the Columbia University Southern millimeter telescope (La Serena, Chile) have a velocity resolution of 1.3 km s" 1 , a rms sensitivity better than 0.12 K, and are vm

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spaced by approximately one beamwidth (8.8'). This is the first survey to encompass the complete latitude and velocity spans of the CO emission from the inner Galaxy. The survey is presented as a collection of 1-V, b-V, and 1-b maps. Several new CO features were observed, and the molecular counterparts of the classic H I structures appear with unprecedented clarity owing to the dense sampling, high sensitivity, and extended latitude coverage of this survey. Features with the largest CO luminosities and velocity widths in the inner Galaxy, outside the nuclear region, were fully mapped and analyzed. The largest of these objects are located at (1, b) = (3.2°, 0.3°), and (5.3°, -0.3°), in a zone previously considered almost devoid of CO emission. Their respective CO luminosities exceed 8% and 3% of the CO luminosity of the inner 500 pc of the Galaxy, integrated over all observed latitudes and velocities. Their velocity widths reach 140 and 100 km s -1 ( FWHM) , respectively. We argue that these features are peculiar molecular clouds located in the vicinity of the Galactic center, and that their large internal kinetic energies (of the order of 10 54 ergs) may have originated in a Seyfert-like event at the Galactic nucleus about one million years ago. Molecular clouds in the surveyed area were identified and catalogued. The inclination and thickness of the molecular layer in the inner Galaxy were measured and their IX

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relevance to current models of the region discussed. A lower limit for the surface density at the center was found, and a CO rotation curve was calculated and compared to available H I rotation curves. x

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CHAPTER I INTRODUCTION The Center of the Galaxy The Galactic center has aroused the interest of astronomers ever since it became clear that the Earth and the Sun are far from the center of the Milky Way, and that a massive galactic nucleus exists in the direction of the constellation of Sagittarius (Shapley 1916 a, b; 1919 a, b, c). The center, however, is hidden from our view by interposing clouds of dust and gas that absorb most of the visible and ultraviolet light emitted in our direction. As a consequence, most studies of the Galactic center have had to rely on radio, infrared. X-ray, and 7 -ray observations. In a paper presented in 1935, reporting radio observations of the sky at 20 MHz, Jansky noted that "radiations are received every time the antenna is directed towards some part of the Milky Way system, the greatest response being obtained when the antenna points towards the center of the system" (Jansky 1935, pg. 1158). This was the first observation of the Galactic center, which also marked the birth of radio astronomy. Following this pioneering work (Jansky 1932, 1933, 1935), the Galactic center region has been observed at many 1

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2 different wavelengths, with increasing sensitivity and resolution, in the continuum and in spectral lines. An enormous volume of data has accumulated. However, because of the complexity of the region, the development of a physical picture that can coherently accommodate all this information has been slow (see reviews by Oort (1977) and Brown and Liszt (1984)). The Galactic center is a singular region. It displays intense activity and characteristics unlike those observed anywhere else in the Galaxy. At the center, near-infrared observations indicate the presence of a spheroidal distribution of late-type stars, in which the density increases inwards (Allen et al. 1983), reaching a maximum near the infrared source IRS 16. Coinciding with the central star cluster and extending some 80 light years in diameter lies Sgr A, a highly complex radio source which is very intense in both thermal and nonthermal emission. Sgr A is also very prominent in molecular line and X-ray emission (Watson et al. 1981), and may be the source of a variable y-ray line produced by positron-electron annihilation at 511 KeV (Lingenfelter and Ramaty 1982). High resolution observations at centimeter wavelengths (Ekers et al. 1983) resolve Sgr A into three components: Sgr A East, Sgr A West, and Sgr A*. The shape and spectral index of Sgr A East indicate that it is probably a supernova remnant. Sgr A West is a strong source of thermal

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3 radiation in which filaments of high-density ionized gas move at radial velocities in excess of 200 km s -1 . If these motions are due to gravitation, the observed speeds would indicate a mass of 3xl0 6 M Q within 0.5 pc of the center (Lacy et al. 1980). At the centroid of Sgr A West lies Sgr A , a nonthermal compact radiosource with a diameter of less than 20 A.U., a spectral index of ~0.25, and a luminosity O A — *1 of 10 ergs s (Lo et al. 1985). Its emission characteristics are best explained, according to Lo (1986), by low level accretion onto a massive (~10 6 M ) Black Hole. o A Black Hole at the Galactic nucleus was proposed by LyndenBell and Rees (1971) to be the source of energetic phenomena in the central region of the Galaxy. At a larger scale, the most striking feature is the radio continuum source known as the "arc" ( Yusef-Zadeh, Morris, and Chance 1984) . Formed by several straight parallel filaments about 100 ly long and 3 ly wide, it is oriented perpendicular to the Galactic plane and apparently connected to Sgr A by curved filaments (see Fig. 1, YusefZadeh et al. 1984). The "arc," which has a spectrum with thermal and nonthermal components (Yusef-Zadeh et al. 1986), may be related to large scale magnetic fields in the center region. Within a few hundred parsecs of the center, there is a large concentration of massive molecular clouds with high velocity dispersions. Observations of the hydroxyl radical

PAGE 14

4 (OH) (McGee et al. 1970, Cohen and Dent 1983), formaldehyde (H 2 CO) ( Scoville, Solomon, and Thaddeus 1972; Cohen and Few 1981) , ammonia (NH 3 ) (Kaifu et al . 1975), carbon monoxide (CO) (Bania 1980, 1977; Liszt and Burton 1978; Heiligman 1982) , and other molecular tracers, indicate that these dense clouds are rich in molecular material; their main component presumably being molecular hydrogen (H 2 ). The distribution of the molecular gas is strongly asymmetrical in both longitude and velocity; most of the emission occurs at positive velocities and longitudes. Some of the clouds in this region apparently form a rotating and expanding ring (Kaifu, Kato, Iguchi 1972) with a radius of 190 pc, expanding at an average velocity of 150 km s" 1 and rotating with a velocity of 65 km s -1 (Bania 1980). Superimposed on this emission, a relatively symmetrical structure is observed in H I and in CO. This feature has two highvelocity wings of emission located in opposite quadrants about the center (see Figs. 1 1-8 and 1 1-9), which were interpreted as the signature of a gaseous rotating "nuclear disk" (Rougoor and Oort 1960). The molecular counterpart of Sgr A seems to be a very intense, overlying molecular cloud with a mean velocity of 40 km s ^ . This " + 40 km s ^ cloud" contains the most intense CO lines in the central region and appears to be connected by a bridge of CO emission to another remarkable molecular cloud of the central region, the Sgr B2 complex.

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5 This unique object contains several compact H II regions and is extremely rich in complex molecules, owing to the combination of large mass with high density. Most of the known interstellar molecules were first observed here. Peculiar phenomena in the central region are found even at a larger scale. Vast gaseous features are seen engaged in large noncircular motions up to 4 Kpc from the Galactic Center (Oort 1977, Rougoor and Oort 1960). The most conspicuous of these is the "3 Kpc arm;" a large structure spanning galactic longitudes from about +12° to -22°, where it apparently reaches a tangential point. It is observed in absorption in the direction of the Galactic nucleus at -52 km s _1 in both H I (Van Worden et al. 1957) and CO (Bania 1980). This feature is thought to be a spiral arm located some 3 kpc in front of the nucleus, and expanding away from it. Another large feature extending across 1=0° with a noncircular velocity is the " + 135 km s" 1 arm." It is probably located behind the nucleus, since no absorption has been observed, and it also seems to move away from the center (see Figs. 1 1-8 and 1 1-9). Together, the "3 Kpc arm" and the "+135 km s“^ arm," constitute the largest deviation from circular motion in the Galaxy, each containing a mass of order 2 x 10^ M Q , and a total expansion energy of 4 x 10^ ergs (Bania 1980). The interpretation of these large, apparently expanding, gaseous structures is quite complex. Attempts to

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6 explain them include expulsion from the nucleus (Van der Kruit 1970, Cohen and Davies 1976), a rotating and expanding tilted disk (Liszt and Burton 1978), elliptical streaming motions (Peters 1975), gas in dispersion orbits (Simonson and Mader 1973), an eccentric bar-like structure (Burton and Liszt 1980), an oval distortion of the central region of the Galaxy (Yuan 1985), and gas moving in a triaxial potential (Heiligman 1982, Vietri 1986, Mulder and Liem 1986). In each case there are features not satisfactorily explained by the models, and the underlying dynamical mechanisms in some of them are not clear. Thirty years after their discovery, the "expanding features" remain some of the most puzzling characteristics of the central region of the Galaxy. In summary, the Galactic center is "sui generis," a region with unusual behaviour observed in scales ranging from a few astronomical units to several kiloparsecs. It displays some characteristics similar to those of active galactic nuclei, although at a much lower level. The study of the Galactic center is not only fundamental to the understanding of the structure and dynamics of our own Galaxy, but may also provide important clues on the phenomenon of nuclear activity in spiral galaxies.

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7 CO Observations of the Galactic Center Observations of the interstellar material in the central region of the Galaxy are the main source of information on the mass density and the physical conditions prevalent at the center; an important fraction of our knowledge about this region has come from observations of interstellar molecules. Although H 2 is the most abundant interstellar molecule, it is very hard to observe directly from the ground. H 2 is an homoatomic molecule, therefore, it lacks a dipole moment and dipole transitions. Near IR quadrupole rotationvibrational transitions can be observed from the ground, but this radiation arises only in quite hot regions. The easiest way to study the cold molecular gas in the Galaxy is to observe the dipole rotational transitions of the second most abundant molecule in molecular clouds, CO. Collisions with H 2 molecules readily excite rotational states of the CO molecules. The first rotational state of the CO molecule, corresponding to an excitation temperature T„ = 5.5 K, is thermalized (T x = T^) by collisions at H 2 densities of a few hundreds per cm^ . The fundamental transition J=l-»0 is readily observed at 2.6 mm (115.2 GHz). This transition is an acknowledged primary tracer of H 2 in molecular clouds (Thaddeus 1977, Dame 1984, Scoville and Sanders 1986). Most of the CO observations of the Galactic center to date were obtained from the Northern Hemisphere with

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8 instruments of relatively high angular resolution (~1'). As a result of these constraints, observations of the latitude distribution of molecular gas within a few kiloparsecs of the Galactic center are moderately to severely undersampled. The better sampled observations were restricted to small areas around the nucleus and near the Galactic plane (Inatani 1982; McCutcheon, Robinson, and Whiteoak 1981; Heiligman 1982). At a larger scale, however, the out-ofplane observations consisted of longitude strips at a few latitudes, with observations typically spaced by 0.5° (~30 beamwidths) (Liszt and Burton 1978; Bania 1980; Sanders, Solomon, and Scoville 1984, Bania 1986). Grids this coarse can easily miss molecular clouds located near the center. The results of the out-of-plane survey of the first galactic quadrant by Cohen et al. (1980), emphasize that well-sampled observations (sampling interval ~1 beamwidth) are necessary in order to detect the continuity of largescale features. These considerations suggest that new, important information on the Galactic center region can be gained from a large scale, well-sampled survey of the latitude distribution of CO emission. Present Work On the premise that out-of-plane features are allimportant for the understanding of the kinematics of the

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9 central region of the Galaxy and impose critical constraints on models, the main goal of this dissertation is to make a sensitive, homogeneous, and well-sampled survey of the latitude distribution of CO emission in the inner 4 kpc of the Galaxy covering all significant out-of-plane and highvelocity emission. For this purpose, we have used the Columbia Southern millimeter-wave telescope at Cerro Tololo Interamerican Observatory, in Chile, to map an area of 100 sq. deg. with an 8.8' beam. The observations are spaced by 0.85 beam within 1° of the Galactic plane, where most of the CO emission is concentrated, and by 1.7 beams for 1° < | b | < 2°. The combined advantages of observing from the Southern Hemisphere and of using a telescope with a relatively large beam, have allowed us to significantly increase the sampling density, improve the sensitivity and extend the spatial coverage with respect to previous observations. The main questions that this work addresses are 1. What is the large scale, two-dimensional (l,b) distribution of molecular gas in the inner 4 kpc of our Galaxy? 2. What are the physical characteristics of the molecular gas in the center region? 3. Are there any significant out-of-plane features missed in previous observations?

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10 4. Is the paucity of molecular gas observed outside the nuclear disk partially caused by lack of latitude information? 5. Do these observations impose new constraints on current models of the center region? This dissertation is organized in six chapters and four appendices. Chapter II describes the Columbia Telescope, the observational techniques used, and presents the large scale characteristics of the molecular gas. Out-of-plane features of great CO luminosities and velocity widths, fully sampled for the first time in this survey, are discussed in Chapter III. Chapter IV presents a catalogue of molecular clouds in the central region, and their identification with H II regions when possible. The kinematics of the molecular gas at the Galactic center are discussed in Chapter V. In this chapter current models of the region and their compatibility with present data are also discussed. A general summary is presented in Chapter VI, and the survey maps are presented in the appendices.

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CHAPTER II THE SURVEY Instrumentation The Galactic center CO survey was carried out with the Columbia Southern Millimeter-wave Telescope (Cohen 1983; Bronfman et al. 1986), located at Cerro Tololo InterAmerican Observatory (La Serena, Chile). Figure II-l is a diagram of the telescope, and a summary of its main specifications is found in Table II-l. This Cassegrain Telescope, a fairly close copy of the Columbia Millimeter-wave Telescope in New York City (Cohen 1977), is an ideal survey instrument. At 115 GHz, the frequency of the CO J=l-*0 rotational transition, its 1.2 m aperture gives a beam 8.8' wide (FWHM). This beam size allows dense sampling of a large area of the sky in a practicable period of time, yet it is small enough to resolve a typical giant molecular cloud located over 20 kpc away (the linear resolution at the Galactic center (Ro=8.5 Kpc) is 18.5 pc) . Computer control of the altitudeazimuth mount and the dome allows fast pointing changes: switching between positions five degrees apart takes less than one second. This allows the use of reference positions located several 11

PAGE 22

12 degrees away from the Galactic Equator with little idle time . The telescope is equipped with a sensitive and stable superheterodyne receiver which has a single-sideband noise temperature of 380 K. The receiver's first stage, consisting of a Schottky barrier diode mixer and a GaAs FET amplifier at the IF of 1390 MHz, is cooled to 77 K by liquid nitrogen, while the second stage of amplification is kept at room temperature. Owing to the low noise temperature of the system and the good atmospheric conditions at Cerro Tololo, integration times are typically quite short; a rms noise temperature uncertainty of 0.1 K is generally reached in less than 10 minutes. The spectrometer is a 256-channel filter bank of standard design. Each filter, 0.5 MHz wide, affords a velocity resolution of 1.3 km s _1 at 115 GHz. The total spectral range of the spectrometer is 333 km s ^, adequate to cover most of the Galactic CO emission, except within a few degrees of the Galactic center, where the CO emission reaches velocity widths up to 500 km s -1 . In this region, two spectra centered at different velocities were taken at each position, and combined as explained in the next section . The atmospheric conditions at Cerro Tololo are extremely favorable for millimeter-wave observations. The water vapour opacity per air mass is usually less than 0.1,

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13 allowing observations virtually year around. An added bonus in the case of this survey is the favourable position of the Galactic center as seen from Cerro Tololo. The center transits almost overhead; so most of the spectra were taken through an air mass of about unity. The system was calibrated daily and checked against reference CO sources before the start of the observations. A standard blackbody chopper-wheel technique (Kutner and Ulich 1981) was used to produce antenna temperatures, T A *, corrected for atmospheric attenuation, resistive losses, scattering, and rearward spillover. To produce radiation temperatures, T R , the antenna temperatures should be divided by the telescope's main beam efficiency, n=0.82. A complete account of the calibration procedure and results is given by Bronfman (1986). The pointing of the telescope was checked twice a day by observing the sun. Right ascension and declination scans across the sun's limb were used to locate the center of the solar 115 GHz continuum emission disk to within 15". The long-term pointing accuracy was checked every few months by observing star positions with a small optical telescope which is collimated with the radio telescope (see Fig. 1 1 1 ) This procedure ensured that long-term pointing errors were smaller than 1' (for details see Grabelsky 1985) .

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14 A Nova 4/x minicomputer with 128 kilobytes of memory was used to control the positions of the mount and dome, the frequency of the local oscillator, the data acquisition, and to perform on-line data processing. The computer system and telescope control are described in detail by Grabelsky (1985) . Observations The almost 8000 spectra comprised in the present survey were taken between February and November of 1984. The area observed, a strip of 25° x 4° along the Galactic plane, was sampled every 7.5' (0.85 beam) in both Galactic coordinates within 1° of the equator and every 15' (1.7 beam) beyond. Relevant parameters of the survey are listed in Table I 1-2. Position switching was used for all observations in order to subtract instrumental and sky background. This was done by pointing the telescope alternatively to a target ("on") position and to a CO-free reference ("off") position every 15 seconds and then subtracting the spectrum of the "off" position from that of the target position. The reference positions, selected from obscuration-free regions found in the ESO plates, were verified to be free of CO emission greater than 0.04 K, which is about one-third of the rms noise per channel in the final spectra. Table I 1-3 lists the CO-free ("off") reference positions used in the

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15 survey. Position switching against clean "offs" yields very flat base-lines; so only first order base-lines were subtracted from the spectra. The velocity centroid of CO emission varies strongly with Galactic longitude, changing sign at 1=0°. Full velocity coverage required setting the spectrometer to different central velocities in the appropriate longitude ranges (Table I I — 2 ) . Nevertheless, within a few degrees of the Galactic center, the velocity width of the CO emission exceeded the spectral range of the spectrometer. This problem was circumvented by taking two spectra, centered at velocities 306 km s ^ apart, at each position within 5° of the Galactic center. The two spectra cover a combined spectral range of 640 km s -1 with an overlap of 26 km s -1 (20 channels); this allows ample emission-free sections to set baselines, and enough overlap to check that the spectra match properly. The values of antenna temperature in the overlapping channels were checked to coincide within the rms temperature uncertainty of the spectra. Poor matchings were rejected and reobserved. A combined spectrum obtained with the procedure described above is shown in Fig. I I -2. The high signal-to-noise ratio and flat baseline of this spectrum are characteristic of the whole survey, as can be deduced from the representative spectra shown in Fig. 1 1-3. An important goal of this survey was to cover the full latitude extent of the Galactic CO emission. We found most

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16 of the emission to be confined between b=-l° and b=l°, and therefore our latitude coverage (-2° < b < 2°) seems adequate. However, since H I emission is observed at higher latitudes in opposite Galactic quadrants (at b>0° for 1<0°, and at b<0° for 1>0°), a search was conducted between b=2° and b=2 . 5° for 1<0°, and between b=-2° and b=-2.5° for 1>0°. Only local emission (V~0 km s -1 ) was detected. The basic datum of the survey is the antenna temperature observed from a particular direction in the sky and at a given radial velocity, Tfl^ b j , V k ) . Twodimensional representations of this four-dimensional data array were accomplished by projecting the integral of the temperature over one of the independent variables onto the plane of the other two. Contour diagrams representing integrated temperature in the longitude-latitude (l,b), longitude-velocity (1,V), and latitude-velocity (b,V) planes were used to exhibit different characteristics of the CO emission (see appendices B, C, and D respectively). In the l,b diagrams, velocity-integrated temperature is displayed in the plane of the sky. This representation is useful to compare the CO distribution with that of other tracers that offer no velocity information such as the infrared or radiocontinuum. Kinematic characteristics of the molecular gas are best displayed in 1,V diagrams of latitude-integrated temperature, and the out-of-plane gas distribution is best seen in b,V diagrams of longitude-integrated temperature.

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17 Large-scale Characteristics of the CO Emission This section describes the general characteristics and the main features of the CO emission observed in the present survey. The distribution of molecular hydrogen mass in the plane of the sky is outlined in maps of velocity-integrated CO antenna temperature (W co = j T dV ) in Galactic coordinates (Fig. II-4). Longitude profiles of CO emission ( 1 ( 1 ) = JjTdbdV ), integrated over different latitude ranges, are displayed in Figs. II-5, 6, 7, and the main kinematic features observed are identified in 1-V diagrams integrated over all the observed latitudes (Figs. II-8, 9). The spatial distribution of CO emission in the inner 3.6 kpc of the Galaxy is displayed in Fig. I 1-4. To examine the full dynamic range of the observations, 1-b maps with three different contour level values (approximately 4, 10, and 20 standard deviations, respectively) are shown in Fig. I 1-4. In the three maps of Fig. I 1-4, the same CO emission data were integrated over the full velocity range (-320 < V < 320 km s ^). As can be seen in Fig. II-4, the CO emission is distributed closely along the Galactic plane with the highest intensity concentrated between 1=-1.5° and 1=1-8° and within approximately 0.6° of the Galactic plane. This central source encompasses the large molecular complexes Sgr A and Sgr B, the nuclear disk, and the expanding molecular ring, which are more clearly demarcated in 1-V diagrams (Figs. II-8, 9).

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18 In longitude profiles (Fig. II-5, 6, 7) the strongly peaked central source appears very asymmetrical with most of the emission found at positive longitudes, resulting in an intensity-weighted mean longitude of 1=0.4°. This central region can be resolved into six prominent peaks (Fig. I 1-5), with the most intense maximum located at 1=1.2°. The central source is also asymmetric in Galactic latitude with most of the CO emission originating below the IAU Galactic plane (see Figs. I 1-6 and I 1-7 integrated below and above the Galactic plane, respectively), yielding an intensityweighted mean latitude of b=-0.05°. The total H 2 mass contained in the central region (1 x b = 3.3° x 1.3° = 490 x 190 pc), calculated from the CO luminosity (see Appendix A) under the assumption that all the emitting gas is at the Galactic center (Ro= 8.5 kpc), is Q 2.7 x 10 M q . This value is in agreement, within a factor of approximately 2, with recent determinations of H 2 mass at the center (see Table II-4). It should be noted, however, that all the masses quoted in Table 1 1-4 might have been overestimated by a factor of 10, since the conversion factor from CO luminosity to H 2 column density in the center may be an order of magnitude smaller than that in the disk of the Galaxy (for details see Appendix A). Most of the CO emission of the inner Galaxy is contained in a band approximately 2° wide in latitude, astride the Galactic plane, and extending through all

PAGE 29

19 observed longitudes. The emission occurs preferentially above the plane for negative longitudes and below the plane for positive longitudes. This asymmetry, also observed in H I, results from a tilt of the gas distribution with respect to the Galactic plane. The magnitude of the tilt, however, is controversial (see Chapter V). The most intense CO emission of the inner Galaxy, outside the central source, originates at two features located at 1~3°, b>0° and 1~5°, b<0° (Fig. II-4). In longitude profiles (Figs. II-5, 6, 7) they can be seen as two large peaks at 1=3.2° and 1=5.3°, which correspond to projected Galactocentric distances of 477 pc and 786 pc, respectively. These sources are remarkable, not only for their high intensities, but also for their unusually large velocity widths (Fig. 1 1 — 8 ) . Their combined luminosity reaches over 1/10 of the intensity of the central source and they are seen in a zone previously considered "almost devoid of CO emission" (Bania 1986, pg. 873). The full extent and intensity of these features were missed in other CO surveys of this region owing to undersampling. The characteristics and origin of these "wide-line" features are discussed extensively in Chapter III. Other discernible peaks outside the central region observed in Fig. II-6 at l=-3.8°, -4.4°, and -5.3° correspond to smaller wide-line features (see Table III-l). The main peak in the central source at 1=1.2°, b>0°

PAGE 30

20 (Fig. II — 7) appears remarkably similar in shape and intensity to the wide-line feature at 1=3.2°. This similarity may suggest that this peak is another wide-line feature concealed by the complex emission of the central region. The kinematics of the CO emission in the inner Galaxy are characterized by asymmetry and noncircular motions. The asymmetry in velocity can be clearly seen in Fig. 1 1-8, where most of the emission occurs at positive velocities with an intensity-weighted mean velocity of 23 km s”^ for the central source. If all the molecular gas in the Galaxy were engaged in circular differential rotation around the Galactic center, the CO emission displayed in the 1-V plane would take place only in the quadrants where velocity and longitude have equal signs. However, large features are also seen in the other two quadrants (Fig. 1 1 — 8 ) indicating the presence of gas at noncircular, "forbidden," velocities. These deviations from circular motion were discovered in the first H I surveys of the center. Figure 1 1-8 displays the kinematic arrangement of molecular gas in the inner-Galaxy in the 1-V plane, integrated over all observed latitudes. Owing to its dense sampling, high sensitivity, and extended latitude coverage, the present survey has revealed new CO features and the molecular counterparts of classic H I features with unprecedented clarity. The following is a brief description of the CO features observed in the 1-V plane, in order of

PAGE 31

21 their relative intensities. To aid in the identification, the main features observed in Fig. 1 1-8 have also been represented schematically in Fig. 1 1-9. The names of the H I features used by Van der Kruit (1970) are given in parenthesis when appropiate. 1. The main maximum. An intense ridge of emission seen along all observed longitudes at low radial velocities (V ~0 km s“^). Its large latitude extent and low velocity (see Appendix C) suggest that the emission originates mostly from local material. 2. The nuclear disk ( VdK IV). A high-velocity wing of emission that extends in the plane from 1=-1.5°, V=-230 km s" 1 towards the center and reappears at positive longitudes and velocities reaching 1=1. 5°, V=220 km s ^ (Fig. II-8, 9). This high-velocity feature is symmetric about the center and was interpreted as a gaseous disk in circular rotation about the center by Rougoor and Oort (1960). 3. The 3 kpc arm . A long ridge of emission extending from 1=13°, V~0 km s ^ across the whole 1-V diagram to 1=-12°, V=130 km s" 1 (Figs. II-8, 9). Crossing 1=0° the 3 kpc arm is seen in absorption against Sgr A at V=-52 km s ^; therefore it is located in front of the nucleus and appears to expand away from it. The 3 kpc arm is the brightest H I feature displaying noncircular motion; it extends to l=-22° where it

PAGE 32

22 apparently reaches a tangent point (Cohen 1975). The 3 kpc arm has been modeled as a ring sector (Bania 1980) and as waves excited at the resonance by an oval distortion of the central region (Yuan 1985). No signs of star formation have been observed in it (Lockman 1980). See, however, chapter IV. 4. The 135 km s— expanding arm ( VdK I_) . A large structure extending from 1=13° to l=-5° at positive velocities. It crosses 1=0° at 135 km s _1 and no absorption is observed. Therefore, it is probably located beyond the nucleus and apparently expanding away from it. It has been modeled as a ring sector by Bania ( 1980) . 5 The molecular ring . This ring of molecular clouds at the Galactic center (not to be confused with the molecular ring at R= 4 to 6 kpc), is a large curved structure mostly located at negative velocities near the center, partially superimposed on the nuclear disk. It extends from 1=-1.2°, V~0 km s -1 to l=-0.5°, V~140 km s 1 and curves back towards the V=0 km s ^ at 1=1.5°. The molecular ring extends below the plane to b=-0.5° (Figs. C-17,19). This feature has been modeled as an expanding and rotating molecular ring by Scoville et al. (1972) and Bania (1980). The prominence of this feature in formaldehyde absorption (Bieging et al. 1980)

PAGE 33

23 indicates that it is probably located in front of the nucleus . 6. Sgr A. This prominent molecular cloud at V~40 km s -1 has been identified with the Galactic center because of its location and high intensity. A deep absorption feature may indicate that this cloud is partially in front of Sgr A*. 7. Sgr B and B2 . Correspond to the large molecular complex at 1~0.5°, V~60 km s _1 . One of the most intense sources observed contains several compact H II regions, complex molecules, and is the site of active star formation. 8. The central gas layer . An intense bridge of CO emission, appearing to connect Sgr A and Sgr B along the Galactic plane at positive velocities (see Figs. C-15, 17). This apparently continuous layer is a unique object in the Galaxy because of its smooth space distribution and large velocity width. 9. The connecting arm ( VdK Ilia ) . Rougoor (1964) called this feature the connecting arm because it seems to connect the nuclear disk, at 1=2.5°, V=270 km s -1 to the main maximum around 1=13°. The connecting arm joins the nuclear disk below the plane at b=-0.25° and, as it approaches the main maximum, its velocity and mean latitude diminish. It can be clearly identified in the 1-V diagrams of Appendix C (Figs.

PAGE 34

24 C-17 to 25) and in the 1-b maps of Appendix B (Figs. B1 to 5) . 10. The 165 km s — feature ( VdK XV). Although discovered in H I, this feature is more prominent in CO. It crosses 1=0° at V~165 km s ^ in the plane; since no absorption against the nucleus is observed, this feature is probably located behind the nucleus apparently moving away from it. It has been proposed to form part of the molecular ring (Scoville et al. 1972) . 11. Norma and Scutum arms ? Probably correspond to the features seen betwen the 3 kpc arm and the main maximum, at negative longitudes. They appear similar to the 3 kpc arm in longitudinal extension and slope in the 1-V plane. Although fairly confused, two structures can be identified in this region and followed towards the center. Absorption features at V=-30 km s — 1 and V=-10 (Fig. D-28) indicate that these structures probably lie in the near side of the Galaxy. A string of clouds observed between l=-2° and 1=2° at V~ 15 km s" 1 at b=-0 . 38° (Fig. C-18) has been identified by Heiligman (1982) with the ScutumCrux arm. The string of clouds between l=-4° and 1=-12° and about V=-30 km s -1 could be the continuation of the arm, while the clouds in the same region at about V=60 km s" 1 , may belong to the Norma arm.

PAGE 35

25 12. Van der Kruit XIV . A symmetric counterpart to the connecting arm at negative longitudes and velocities. It can be seen in Figs. C-12 to C-15 at l=-2° to -5° and between V~ -220 to -150 km s _1 It joins the disk at b=0° and continues above the plane towards the main maximum. It is much weaker and less well defined than the connecting arm. However, this is one of the few features displaying symmetry about the center . 13. Van der Kruit I I . Also known as the 70 km s _1 arm, is a faint and narrow feature in the plane. It has been studied in detail by Shane (1972). 14. Feature I I lb . An in-the-plane feature which runs parallel to the 70 km s' 1 arm and seems to end near the wide-line feature at 1=5.3°. 15. Van der Kruit IX . Found at higher velocities than feature Illb, crosses the connecting arm and becomes confused with emission of near-by features near 1=7° 16. Van der Kruit XI I . A weak cloud detected at 1=4°, b=-2.4°, and forbidden velocity V=-83 km s" 1 , seems to be the first CO detection of feature VdK XII, the most prominent high-velocity feature seen in H I (Cohen 1975). A nearby cloud at V~30 km s“^ coincides with the 1-V locus of the 3 kpc arm but lies well below the plane ( b=-2.2°). Not shown in Fig. I 1-9 .

PAGE 36

26 17. A feature extending below the plane between 1=3.3° and 4.0° and between velocities V=40 and 150 km s _1 was detected for the first time in this survey. It could be related to H I feature J4 (Cohen 1975) but their slopes in the 1-V plane are different (Fig. D-19). Not shown in Fig. II-9. The characteristics of the molecular clouds that form many of these features are tabulated in a catalogue in Chapter IV.

PAGE 37

TABLE 2.1 TELESCOPE CHARACTERISTICS Aperture 1.2 m Beam width (FWHM) 8.8' at 115 GHz Effective F/D 3.79 Beam Efficiency 0.82 Noise Temperature 380 K Single sideband Spectrometer 256 channels 0.5 MHz wide

PAGE 38

TABLE 2.2 PARAMETERS OF THE SURVEY Angular resolution 8.8' Velocity resolution. ... 1 . 3 km/s Sensitivity 0.12 K rms per velocity element Spatial coverage: Galactic longitude 348° to 13° Galactic latitude -2.0° to 2.0° Sampling interval: o o o A [b] o o rH V 0.125° = 0.85 beam M o o A [b] 0 o CS) V . r. .0.250° = 1.70 beam Velocity coverage: 348. 0°< 1 <353.5° -244 to 00 00 km/s 353 .5°< 1 <355.5° -179 to 153 km/s 355. 5°< 1 < 5.0° -319 to 319 km/s V o o Li") 1 < 13.0° -88 to 244 km/s

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29 TABLE 2.3 CO-FREE REFERENCE POSITIONS 1 b 1 b (°) (°) (°) (°) 347.00 3.00 0.70 -3.30 348.00 -5.00 1.00 -3.90 348.00 4.00 1.90 -5.35 352.00 -5.00 3.80 -4.80 353.00 -5.00 4.75 2.50 353.00 4.00 5.00 -4.00 354.00 -5.00 7.00 5.00 355.00 5.00 8.00 -5.00 355.00 5.00 8.50 2.75 356.04 -5.23 9.00 5.00 357.00 -5.00 10.00 5.00 357.00 3.00 11.00 -4.00 359.80 -5.75 12.00 5.00 0.00 -4.00 14.00 2.50 0.00 -3.00 15.00 -5.00

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30 TABLE 2.4 MOLECULAR MASS IN THE GALACTIC CENTER Reference Observations Area lxb pc Mass 10 8 M q Audouze 1979 12 CO / 13 C0 / C 18 0 600x600 1.3 Heiligman 1982 13 CO 340x140 1.2 Sanders et al. 1984 12 CO 3000x320 3-5 Bania 1986 12 CO, 13 CO 550x90 3.3 This work 1987 12 CO 490x190 2.7

PAGE 41

31 Fig. II-l Diagram of the Columbia Southern millimeter-wave Telescope.

PAGE 42

12 CO spectrum in the direction of the Galactic center (l / b=0° < 0°) formed by combining two overlaping spectra, each 333 km s _1 wide. Fig. I 1-2

PAGE 43

33 2.000 1.750 1.500 1.250 1.000 0.875 0.750 0.625 0.500 0.375 0.250 0.125 0.0 -0.125 -0.250 -0.375 -0.500 -0.625 -0.750 -0.875 1.000 -1.250 l • 0.125 t I I T I I I I I I I I I I I I I I I 1 M I I I I I I I 1 ^ Aa/\jL JL 0.0 i r m -rr' n M I I i i i i i i i i i i r r n i i i i -1.500 -1.750 2.000 I 1 I I 1 1 I I I 1 1 1 8 8 8 8 8 8 o LSR RfiOIAL VELOCITY (KM S'' ) 8 3 8 Fig. 1 1-3 Representative survey spectra; the baselines are spaced by 6 K.

PAGE 44

(0 (0 C P C d a) p i 01 •H T 3 m >1 >1 G m g 0) IT} P Own} U OH T 3 XII O Eh B P 1 * CTAJ d 01 O P CM G 1 c P p d O o 0) CM G m d i oi P IT} a) it} > A a) iH 01 g o p -p o -P XI ai X3 rH 01 T> T 3 HI T) G > -P g •H P 0 ai a 01 a ai a* p d o o ai > 04 o -p Q) jg p p c Q) P -P ai ip p on pi 01 G . . O 0 -P OjP 01 (T} ai -p p > p ai 0 id a T 3 01 p •P d A T 3 Eh G (TJ • P >1 01 01 o CM •P Q) 4H ^ p ^ G X-P G O G O ai -d a} m ai d •P p ^ ID a P 3 0 )P 01 01 «. 3 P P d

d P ai d> ai p * ^ P a 1 ai in 01 01 g p d p 0 a •P a) G rP c CN d 0 a P -P Oj *P co g p H M d> •p &4

PAGE 45

35 CM e O CM I GALACTIC LONGITUDE

PAGE 46

7000 Fig. II-5 Longitude profile (double integral) integrated between V=-320 and 320 km s -1 , and between b=-l° and 1°.

PAGE 47

37 Fig. I 1-6 Longitude profile integrated between V=-320 and 320 km s“ , and from b=1° to 0°.

PAGE 48

Fig. II-7 Longitude profile integrated between V=-320 and 320 km s 1 , and from b=0® to 1°.

PAGE 49

39 13.0 8.0 3.0 2.0 7.0 12.0 i i i i i i i i i i i i i i i o i\j .e a co — — — — — ro U)r\jroroi'jr\> — — — — — a> cr> .c i\j • ooooor\j.ccnoo ooai^njoooCTicrooooooo* • • • aooooo ooooooaoooo• • oooo oooo aooooo ooooooaaooo Fig. I 1-8 L-V diagram integrated between latitudes b=-2.5° and 2.5°. Contours are set at 2 K from 2 K. 300.0 280.0 260.0 240.0 220.0

PAGE 50

ooooooooooo Fig. I 1-9 Squematic 1-V diagram, the numbered structures are identified in the text.

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CHAPTER III WIDE-LINE CLOUDS NEAR THE GALACTIC CENTER The most remarkable objects mapped in our survey are the wide-line features located at (l,b) = (3.2°, 0.3°) and (5.3°, -0.3°). Their respective CO luminosities exceed 8% and 3% of the CO emission of the inner 500 pc of the Galaxy, integrated over all observed latitudes and velocities (see Fig. II-4). Thus, the wide-line objects are the most intense localized CO sources in the Galaxy beyond the nuclear region. The velocity spans of these outstanding features are the greatest observed for molecular gas in the Galaxy outside the nuclear disk. The velocity widths ( FWHM) reach 140 km s -1 for the object at 1= 3.2° and 100 km s -1 for the object at 1= 5.3° Their CO emission is remarkably smooth along the large velocity spans, even though some structure is also apparent (Fig. III-l). The unusual characteristics of these objects are particularly notable in 1-V diagrams (Fig. III-2), integrated over 1° in latitude below (a) and above the Galactic plane (b) . The wideline objects appear as intense horizontal emission streaks at 1= 3.2° and 5.3°. Contrary to their large velocity extents, their longitude widths are 41

PAGE 52

42 small, confined to about 0.5° (4 beamwidths). This apparent "longitude-crowding" morphology that the wideline objects present in the 1-V plane contrasts sharply with the 1-V signature of known arm-like structures, which span up to tens of degrees in Galactic longitude and usually have velocity widths well under 40 km s' 1 . Although previous observations of the central region of the Galaxy in OH (McGee et al. 1970) and CO (Bania 1977) have offered some evidence of peculiar "clumps" or "features" with surprisingly wide lines, large portions of their emission were missed because of undersampling. Owing to the improved sampling density and latitude coverage of our survey we have been able to determine, for the first time, the total extent and CO luminosity of these objects, thus revealing their real importance (Bitran et al. 1985). Other, less prominent wide-line features can be seen at (1, b, V) = (-5.3°, 0.4°, 84 km s' 1 ); (-4.4°, 0.6°, 72 km ' s -1 ); and at (-3.8°, 0.9°, -83 km s' 1 ) (Fig. III-2). It is possible that the objects at 1= 3.2° and 5.3° are the largest members of a hierarchy of wide-line clouds in the inner Galaxy, some of which may even be concealed by the complex emission of the nuclear region, for example at (1, b ) = (1.2°, 0.3°) (Fig. II-4).

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43 Kinematic Features or Molecular Clouds ? The interpretation of these unique objects poses several difficult questions, and startling possible consequences. The first question that arises when pondering the nature of the wide-line features is whether they are: a) kinematic features-superpositions of many clouds along the line of sight, like the tangent point of a spiral arm— or b) real objects—localized condensations like molecular clouds found elsewhere in the Galaxy, whose structure along the line of sight is similar to that observed in the plane of the sky. The first interpretation has been adopted by Bania et al. (1986) and by Stark and Bania (1986), who have observed two of our wide-line clouds ((l,b) = (3.2°, 0.3°) and (-5.3°, 0.4°)) at high resolution in 12 C0, 13 C0 and CS. Bania et al. (1986) resolve the cloud at (l,b) = (-5.3°, 0.4°) into 3 molecular complexes with smaller velocity widths. They argue that the 3 components are not bound to each other and attribute their clustering in space and velocity to gravitational forces working on a galactic scale. As for the wide-line cloud located at (l,b) = (3.2°, 0.3°), observations by Stark and Bania (1986) resolve the cloud into 16 components. They argue that the 16 components are aligned along the line of sight forming a spiral arm, or dust lane in the inner Galaxy.

PAGE 54

44 Although possible, we consider kinematic explanations for the wide-line clouds extremely unlikely. A kinematic interpretation is most unlikely for a collection of wideline clouds, since it would imply an extremely privileged location for the observer. It is also an unlikely explanation for an individual wide-line object because, in addition of requiring alignment of several clouds along the line of sight, the straight horizontal tracks of the wideline clouds in the 1-V diagrams (Fig. III-2) do not resemble the 1-V shapes of recognized kinematic features, such as, for example, the tangent point of the Carina arm. While the tangent point of the Carina arm is symmetrical about the plane in b, it is highly asymmetrical in 1, with a well defined boundary in one direction (at the arm's edge) and diffuse at the opposite direction (Cohen et al. 1985). In contrast, the wide-line objects have no symmetry about the Galactic plane but are symmetrical in 1 with well defined boundaries in both directions. Almost conclusive evidence against a kinematic interpretation, moreover, is provided by the nearly complete abscence of a 21 cm counterpart to either wide-line cloud (see e.g. Sinha 1979), since kinematic structures such as spiral arms, tangent points, or dust lanes are characteristically conspicuous in H I emission. Based on the above discussion, we conclude that the wide-line features are most likely to be high-density

PAGE 55

45 localized clouds similar to molecular clouds found elsewhere in the Galaxy, albeit peculiar because of their large internal motions. This conclusion, we emphasize, does not preclude internal structure, which is generally characteristic of large molecular clouds and cloud complexes in the Galaxy (see e.g. Dame 1984). Because the wide-line clouds are closely confined to the Galactic plane, cluster about the Galactic center, and have high radial velocities, they are almost certainly a population associated with the Galactic center. Therefore, they must be located not much further from the Galactic nucleus than their projected galactocentric distances, 0.4 to 0.8 kpc . The radii, masses, and energies of internal motions for the wide-line clouds have been calculated on the assumption that they are molecular clouds located at the Galactic center (Ro=8.5 kpc). These values, along with other observational parameters, are listed in Table III-l. The Masses Owing to their high intensity and fairly sharp edges, the wide-line clouds are defined with little ambiguity against the CO background emission of the inner Galaxy. After subtraction of background emission in Wqq maps integrated over the relevant velocity ranges, effective

PAGE 56

46 o radii (r) are defined so that nr is the area within the W co contour at 1/e of peak intensity. Molecular masses, Mqq, are derived from the W co summed over the effective area of each cloud using the conversion factor (Bloemen et al. 1986) , ( N H 2 ) 20 -2-1-1 = 2.8 x 1CT U mol cm K 1 km s (3.1) W CO and a mean molecular weight of 2.76 x 10 -24 gm (Allen 1973). A detailed description of the calculations of the masses is found in Appendix A, and their values are listed in Table III-l. It is important to note, however, that using the above conversion factor, which was derived from gamma-ray and CO observations of the outer Galaxy, may lead to an overestimation of the masses of clouds near the Galactic center by as much as a factor of 10 (see "Molecular Masses near the Galactic Center" in Appendix A). Therefore, Table III-l also lists the lower limit for the masses of the wideline clouds, M co /10. Because of their extraordinary characteristics, the two leading wide-line clouds at 1=3.2° and 1=5.3° were also fully mapped in 13 C0. We found that the clouds' line intensity ratio 13 CO/ 32 CO was approximately 1/10, or about half of the typical Galactic value. Integrated 32 CO and 13 CO spectra, summed over the angular extent of the clouds

PAGE 57

47 are displayed in Fig. III-3 for the cloud at ( 1/ b) = (3 . 2°, 0 . 3°) , and in Fig. III-4 for the cloud at ( 1 / fc>) = ( 5 . 3° , -0 . 3° ) . Using the spatially coincident ^ 2 C0 and 12 CO observations and assuming local thermodynamic equilibrium (LTE), masses for these clouds were calculated following the LTE method outlined in Appendix A. The values for the LTE masses, M LTE = 5 x 10 6 M Q for the 1=3.2° cloud and = 1 x 10^ M Q for the 1=5.3° cloud, are within the range of the masses calculated from 12 C0 luminosities in Table III-l, but closer to the lower limit (M co /10) . The virial masses of the wide-line clouds, calculated on the assumption of spherical symmetry with uniform mass distribution and using the observed velocity dispersions (see virial method in Appendix A), are an order of magnitude larger then the masses derived from the CO luminosities (Table III-l). If the standard value for the conversion factor N(H 2 )/W C q indeed overestimated the CO masses, the discrepancies are even greater. The 3.2° and 5.3° clouds seem to be out of virial equilibrium, apparently expanding almost freely on a time scale t (calculated as the radius divided by the rms velocity width) of 0.7 x 10 6 years. Calculations for the smaller wide-line objects give similar results with time scales of approximately 1 x 10^ yr (Table III-l) .

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48 Origin of the Energy The kinetic energies of internal motions of the two leading wide-line clouds, calculated assuming that they are spherically symmetric, homogeneous, and expanding, range between 10 52 and 10 54 ergs (Table III-l). A source of energy capable of producing about 10 54 ergs in a small region (r~50pc) is required to drive the largest wide-line cloud. Such high energies are unlikely to be produced by supernovae, stellar winds, cloud collisions or tidal forces, which are the main energy sources for the interstellar medium (ISM). Let us consider first stellar energy sources. The total energy output of a supernova outburst is estimated to 53 be ~10 ergs. The kinetic energy of the ejecta is about 10 51 ergs (Chevalier 1977), following the explosion this energy is transferred to the interstellar medium. Stellar winds from 0 stars are also an important source of energy for the ISM; the stellar wind of a massive 0 star (M > 20 M 0 ) can impart, over the lifetime of the star, a mechanical energy to the ISM that is comparable to that of 'a supernova outburst (McCray and Snow 1979). Assuming then, that both a supernova and an 0 star each deposit at the most ~ 10 51 ergs of kinetic energy in the parental molecular cloud, at least 10 to 10 such events are required to drive the largest wide-line cloud.

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49 Starbursts producing large number of 0 stars and supernovae have been shown to account fairly well for the characteristics of active energetic sources such as the nuclei of M82 and NGC253 (Rieke et al. 1980). In the case of the wide-line clouds, however, the copious ionization and far-infrared emission that such massive star formation would produce, are almost completely absent. The main wide-line clouds are inconspicuous in the IRAS and other far-IR surveys of the Galactic center (e.g. Campbell et al. 1985), practically invisible in radiocontinuum emission (Altenhoff et al. 1978), and only a few weak 5 GHz sources lie in their vicinity (Downes et al. 1980). The deficiency in ionization is demonstrated particularly well by a comparison of the weak 5 GHz continuum sources found in the vicinity of the 3.2° cloud (Altenhoff et al. 1978) with W49, the giant H II region at 14 kpc resulting from approximately 10 0 stars. The only continuum sources within 1° of the 3.2° cloud are at least one order of magintude less intense than W49 and are, therefore, presumably the result of only a few 0 stars or supernovae. On the hypothesis of starburst, therefore, the ionization in the vicinity of the 3.2° and 5.3° clouds would be deficient by 2 to 3 orders of magnitude, and the far-IR emission by at least an order of magnitude. We conclude that, because of the lack of ionization and far-IR radiation, O stars and supernovae cannot account for the kinetic energy of the leading wide-line objects.

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50 Out-of-plane neutral hydrogen observations provide evidence of a population of high-velocity clouds (HVC), distributed through all galactic longitudes outside the disk, with velocities in excess of 80 km s -1 (see review by Hulsbosch 1975). In a Galactic fountain model, Bregman (1980) proposed that the HVC condense from a Galactic corona formed by supernova-heated gas from the disk and subsequently fall back onto the disk. A massive cloud falling from the Galactic halo onto the disk, where it collides with molecular clouds, can transfer large amounts of kinetic energy to a small region in the plane. However, the energy transferred to clouds in the disk by a falling HVC would be several orders of magnitude lower than that required to drive the leading wideline clouds, since the masses and velocities of the HVC are too small. Observed velocities of the HVC rarely exceed 200 km s" 1 , and although their masses and distances are poorly known, the most widely accepted view is that the HVC are only a few kpc out of the plane of the Galaxy (Hulsbosch 1975); therefore, their masses probably range between 10 2 and 10 4 M Q . Moreover, the large disturbances that collisions of the needed magnitude would cause in the disk are not apparent in our CO data. The possibility that the wide-line profiles were caused by Galactic tidal forces disrupting the clouds was also investigated. A necessary condition for a cloud to be stable against tidal disruption is that its self-gravitation

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51 should be greater than the Galactic tidal forces acting on the cloud. For a spherical cloud at galactocentric distance r o' the minimum average density for which the cloud is gravitationally bound is given by (Mihalas and Routly 1968): 3W p c * 2 It G 1W o “ I <30/dr I j (3.2) where G is the gravitational constant, W = 0/r is the average angular velocity, and d0/dr is obtained from a model for the Galactic rotation curve 0(r) (Burton and Gordon 1978). Expressing the mass in solar masses, and the volume in cubic parsecs, the critical density can be written as -4 M P c = 1.11 x 10 * W Q (W Q |d0/dr| ) [— 2 -] (3.3) ° pc J The critical densities for the wide-line clouds, calculated with the above formula, are listed together with the observed densities in Table I I 1-2. Comparison of the corresponding density values indicate that, if the standard conversion factor between CO luminosity and H 2 column density (Appendix A) is valid in the neighborhood of the Galactic center, the wide-line clouds are stable against tidal disruption. On the other hand, if the CO masses of the clouds near the center have indeed been overestimated by a factor of ten as proposed by Oort (1977), the wide-line clouds would be unstable against tidal disruption. In case

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52 that this holds true, tidal disruption would be expected to result in elongation and/or break-up of the clouds radially towards the Galactic center, i.e., the tidally disrupted clouds would appear elongated in Galactic longitude but not in latitude. However, exactly the opposite is observed in the wide-line clouds; they are more extended in Galactic latitude than in longitude. For these reasons it is unlikely that Galactic tidal disruption plays an important role in the production of the wide-line profiles. Are the Wide-Line Clouds Powered by the Galactic Nucleus ? The inability of conventional energy sources to produce enough energy to power the largest wide-line clouds prompts the search for more exotic possibilities. The central position of the wide-line clouds and the magnitude of the energies involved, point to the Galactic nucleus as a likely energy source. Considering the direct observational evidence now accumulating for energetic phenomena at the Galactic nucleus (Lo 1985) and for resemblance to a Seyfert nucleus (Kassim et al. 1985, 1986), the general explanation of the wide-line clouds that we find most plausible is that they are the large molecular clouds that happened to be in the vicinity of the Galactic nucleus at the time of a Seyfert outburst about one million years ago.

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53 The bolometric luminosity of a Seyfert nucleus ranges from 10 43 to 10 46 erg s _1 (Wilson 1982); a typical value being 3.8 x 10 45 erg s -1 (~10' 1 ' 3 L q ) according to Osterbrock (1984). The time scale of a Seyfert event is estimated from the characteristic velocities (10 3 to 10 4 km s -1 ) and sizes (~1 kpc) of the disturbed central regions to be between 10 5 r and 10 years. A moderately bright Seyfert event with a 11 44 -1 luminosity of ~ 10 L q (7.6 x 10 ergs s ), and lasting an average time of 5 x 10 5 yrs, will deposit in the surrounding ISM a total energy of 1.2 x 10 58 ergs. The largest wide-line cloud subtends an angle of about 10 -3 sr as seen from the Galactic center, and in such Seyfert outburst at the Galactic nucleus would be illuminated by 5 6 ~ 10 ergs of radiant energy. A fraction of 1% of this energy transformed to kinetic energy would account for the observed internal motions. In the Seyfert outburst scenario, the large internal motions of the wide-line clouds would be the result of rapid ionization and heating "in situ" by the UV and X-rays of the flat Seyfert spectrum (and possibly by relativistic particles) followed by isothermal free expansion of the 4 resultant hot (~2xl0 K) , high-pressure fireball through the sonic speed ( ~ 10 km s” 3 ). During the outburst, a distant observer might have classified the cloud as a Seyfert forbidden-line filament. In fact the velocity

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54 widths, sizes, and densities of the forbidden-line filaments observed in Seyfert Galaxies (Osterbrock 1984) are comparable to those of the wide-line clouds. On termination of the outburst, without a sustained source of ionizing photons, the electrons and ions will recombine and become neutral in a short time. Lo (1986) quotes a recombination time of less than 10 years for Sgr A West. The recombination time is inversely proportional to the ionized gas density and, therefore, it is likely to be about a few hundred years for the wide-line clouds. Molecular formation would follow on a time scale difficult to specify (since it would depend on the fate of the dust grains), but presumably well within the last million years. Once set in motion, the dense gas of a large molecular cloud is not readily stopped, as Oort (1977) has emphasized, and the wideline objects would be among the most durable remnants of a Seyfert outburst in a spiral Galaxy such as ours. The 3 kpc and the 135 km s ^ classic expanding arms of the inner Galaxy could be descendants of a previous generation of wide-line objects (the masses and the kinetic energies are about right, if some swept-up material is added), and the difference in age between these arms and the wide-line objects would be a rough measure of the time between successive outburst: about 10 Myr. Although the proposed mechanism can account for the high energies needed to drive the wide-line clouds and could

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55 also explain some of the expanding features farther from the nucleus, it has several weak points: It is, of course, debatable that our Galaxy has had Seyfert-like behavior, with an outburst taking place as recently as a million years ago; the mechanisms for the absorption of energy by the clouds and the efficiency of transformation to kinetic energy are unclear; the velocity widths of the wide-line clouds are a few times smaller than those observed in Seyfert filaments and do not surpass the velocities allowed by the rotation curve of the Galaxy. Also, as a consequence of the outburst, the population of wide-line clouds near the Galactic center would be expected to be larger. This may indeed be the case, however, since wide-line clouds located in the central 3° of the Galaxy are hard to recognize because they would blend with the complex emission of the region. In fact, in 1-V diagrams (Fig. II-8, I I I — 2 ) , the central region appears studded by several horizontal streaks which could be unrecognized wide-line clouds. In summary, the wideline features are among the most puzzling phenomena observed in the central region of the Galaxy. Their characteristics suggest that they are localized high-density molecular clouds situated near the Galactic center. Conventional mechanisms cannot provide the energies required to explain their large internal motions. We propose that the Galactic nucleus is the energy source for the wide-line clouds, which implies that our Galaxy

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56 underwent a Seyfert-like outburst about one million years ago .

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57 TABLE 3.1 WIDE-LINE MOLECULAR CLOUDS Obj ect 1°, b° (a) (b) (c) (d) (e) (f) (g) (h) 5.3, -0.3 95 43 786 31 0.7-6. 4 64 35-350 0.7 3.2, 0.3 104 60 477 43 2.3-23 175 250-2500 0.7 -3. 8, 0.9 -83 19 579 19 0.2-1. 7 8 2-18 1.0 -4. 4, 0.6 72 22 658 20 0.2-1. 7 11 3-25 0.9 -5. 3, 0.4 84 20 787 25 0.5-5. 2 12 6-62 1.2 (a) : cloud mean velocity [km/s] (b) : rms velocity width, AV rms [km/s] (c) : projected distance from Galactic nucleus [pc] (d) : effective radius [pc] (e) : CO mass range [ 10 6 M Q ] (f) : virial mass [10^ M Q ] (g) : internal kinetic energy range [ 10 51 ergs] (h) : characteristic time ( r /AV rms j [ 10 6 yr]

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58 TABLE 3.2 DENSITIES OF THE WIDE-LINE CLOUDS Obj ect 1(°), b(°) (a) (b) (<=) 5.3, -0.3 71 7.1 13.6 3.2, 0.3 96 9.6 33.0 -3.8, 0.9 80 8.0 25.2 -4.4, 0.6 69 6.9 20.1 -5.3, 0.4 108 10.8 13.5 (a) : measured cloud density [M 0 /pc 3 ] (b) : lower density limit [M Q /pc 3 ] (c) : critical density [M Q /pc 3 ]

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59 b 1 . 00 * t = 3.250 1 1 1 1 1 1 mm m , mm A J ^ ^ aA a /\ /V\ — I I 1 I I -200 0 200 V LSR (km s" 1 ) Fig. III-l Spectra taken across the 1=3.2° cloud.

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o O CO (D • P (U T3 H (d ,G 3 > •POM •H P d) • P P W • (00 G G W H U1 -rl O T3 (N -H d M P id P > O M -H o O II CT> > X> -O H G -H PC • Id 5 G |H G £ P • 10 OP •HOMO 0 13 O O O P P i-H O P (fl O -M G 10 U P -M d) P O T3 I • d) ft 3 II — A PJ3PP « •H J3 £ Cn g tr>,G o G o p P m 0 M M -H M MPS (fl CM 1 H H H tn •H tM

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61 1 1 ! 1 1 1 1 | l l l l | I l I I — I — l — i — i — i — r • 1 1 1 1 L_ o — 1 1 1 1 L i i i m o in o tn m rO fO 3anii9Ncn oiiovivo LSR RADIAL VELOCITY (km s"

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Fig. I I 1-3 12 C0 and ^ 2 C0 spectra of the wide-line cloud at (l,b)=(3.2° / 0.3°) / summed over the cloud's angular size.

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63 Fig. III-4 ^CO and CO spectra of the wide-line cloud at ( 1, b)=( 5 . 3°, -0 . 3°) , summed over the cloud's angular size.

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CHAPTER IV A CATALOGUE OF MOLECULAR CLOUDS IN THE INNER GALAXY A distinctive characteristic of the molecular gas in the Galaxy, in contrast with the atomic gas, is its clumpy spatial distribution. While the atomic gas is distributed smoothly, the molecular gas is found concentrated in clouds and cloud complexes. The larger complexes, self-gravitating aggregates with masses of the order of 10 5 M Q and radii of approximately 50 pc, are known as "giant molecular clouds" (GMC). They are recognized as the phase of the Inter Stellar Medium (ISM) where most star formation takes place and are useful tracers of spiral structure (Dame 1984). Because of its clumpy nature the distribution of molecular gas in the Galaxy is best quantified by an inventory or catalogue of molecular clouds. A catalogue that lists the location and physical characteristics of molecular clouds in the inner Galaxy is useful to study spiral structure and star formation in this region, and for comparison with molecular clouds in the outer Galaxy. In this chapter we present a catalogue of molecular clouds in the inner Galaxy. The catalogue lists the mean coordinates, radial velocities, radius, velocity width and apparent CO luminosity of the molecular clouds in the region (see Tables 64

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65 IV1 to 8). This is the only catalogue of molecular clouds in the inner Galaxy, available to date, with a complete latitude coverage. The catalogue, which includes the observed highvelocity clouds, is organized by membership of the clouds in recognized H I structures such as the 3 kpc arm, the 135 km s arm, etc. (see Chapter II). The membership of the clouds in these structures was decided by spatial and velocity coincidence of the clouds with the locus of the H I structures as they appear in a survey by Cohen (1975). In several instances, superpositions of features make the identification of clouds and their membership in H I structures uncertain. Therefore, it is likely that the classification of some of the clouds may have to be revised when higher-resolution observations become available. The clouds at the Galactic center are difficult to identify against the intense, wide-velocity background, therefore they have not been included in the catalogue. The wide-line clouds are catalogued in chapter III. The procedure for determining the boundaries of the clouds was done as follows. The clouds were identified in "clipped" latitude-integrated 1-V diagrams. These diagrams are produced by setting to zero the contributions of channels with antenna temperature lower than a certain threshold value. The molecular clouds are seen distinctly in this type of diagram, when the threshold has been

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66 properly set to eliminate the difusse background emission. Through trial and error, a threshold temperature of 2 K was found generally adequate to produce good contrast in the diagrams, without eliminating identifiable clouds. An integrated spectrum (T vs. V) was then produced for each cloud by adding the spectra within the spatial boundaries of the cloud. The boundaries of the clouds were found by inspection of the 1-V, 1-b, and b-V diagrams displayed in appendices B, C, and D. The limits of the clouds were revised by examining temperature profiles, and the procedure iterated until all the emission from the cloud was included in the integrated spectrum. The location and velocities of the clouds in the catalogue are given as the intensity-weighted mean coordinates. We have chosen these coordinates rather than the more commonly used, and easier to find, peak coordinates because they are more physically meaningful. The intensityweighted coordinates are closer to the center-of-mass coordinates of the cloud, since the CO intensity is proportional to the molecular mass (see Appendix A) The intensity-weighted mean longitude was calculated as. (4.1) similarly the mean latitude is.

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67 < b > (l,b,V) b ( 1, b, V) and the mean velocity. T i ( l,b, V) V E T d,b,V) 1 / b, V 1 (4.2) (4.3) The listed radius corresponds to the equivalent radius of a circle having an area equal to the cloud's observed area, and is calculated as r = (A/rt) 1//2 . The velocity widths of the clouds, found in Tables IV 1-8, correspond to the FWHM calculated from Gaussian analysis of the integrated spectra. The integrated spectra were also used to calculate the apparent CO luminosity, I co . This quantity was calculated according to the expression: I co = Tsum x D1 x Db x DV, where the area under the integrated spectrum is Tsum, the spatial picture elements are Dl=Db=0. 125°, and the velocity resolution element is DV=1 . 3 km s _1 . The apparent CO luminosities are followed by their corresponding rms uncertainties . Because of the large noncircular motions observed in the central region of the Galaxy, kinematic distances to clouds in this area are not reliable; therefore they are not listed in the catalogue. For some clouds, however, it has been possible to find associated H II regions by spatial and

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68 velocity identification in catalogues of H II regions (e.g. Georgelin and Georgelin (1976), Lockman (1979), and Downes et al.(1980)). When an estimate for the distance of the associated H II region was available, the molecular mass of the clouds was calculated. The mass of a cloud can be calculated directly, in solar masses, from the CO luminosity I co , and the distance d in kpc, with the formula M = 1.9 x 10 3 I co d 2 (4.4) For details on the mass calculations see Appendix A. Notes on Selected Clouds . We have examined catalogues of H II regions looking for counterparts to the molecular clouds listed in the catalogue. Associations of molecular clouds and H II regions provide evidence of star formation and of real density enhancements as opposed to features produced by velocity crowding. The identification of molecular clouds with H II regions is far from exhaustive because of incompleteness of the available H II catalogues in spatial and velocity coverage. Cloud # 15, 3 kpc arm. This cloud located at (1, b, V) = (5.35°, -0.03°, -20 km s' 1 ) may be associated to the H II region at (1, b, V) = (5.479°, -0.241°, -26 km s' 1 ) (Downes et al. 1980). This association would provide substantial • evidence of star formation in the 3 kpc arm, contradicting

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69 earlier observations (Lockman 1980) in which no tracers of star formation were found. If this cloud is located at a distance of 5.5 kpc, its molecular mass would be 4.7 x 10~* M oCloud # 3, III b. Cloud # 3 at (1, b, V) = (7.56°, 0.05°, 122 km s' 1 ) is probably associated to the H II region at (1, b, V) = (7.472°, 0.06°, 121 km s' 1 ). Downes et al. (1980), argue that this source is located at 27±5 kpc. This distance implies a molecular mass of 3±1 x 10 6 M Q . Cloud # 8, III b. The location of this cloud, (1, b, V) = (12.17°, -0.02°, 118 km s' 1 ) indicates an association with a source at (1, b, V) = (12.284°, -0.116°, 119 km s' 1 ). The distance in Downes et al. catalogue, 16±1 kpc, implies a mass of 1.2 x 10 6 M Q .

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70 TABLE 4.1 MOLECULAR CLOUDS IN THE 3 KPC ARM # <1> (°) (°) ( km/s ) V r (km/s) (') : co (a) O" (a) 1 -10.24 0.13 -107 10 15 4.18 0.03 2 -9.37 0.43 -105 13 15 6.33 0.04 3 -8.47 -0.29 95 8 25 7.89 0.04 * 4 -7.79 -0.14 93 6 22 7.43 0.04 * 5 -7.46 -0.11 84 10 25 13.25 0.05 * 6 -6 . 66 0.08 87 11 31 21.49 0.07 Je 7 -5.87 0.13 85 9 21 5.24 0.04 8 -4.49 0.19 76 13 16 7.49 0.03 9 -3.29 0.29 70 7 13 4.13 0.02 10 3.52 -0.49 28 11 23 9.43 0.04 11 4.28 0.45 28 9 12 3.21 0.07 12 4.76 0.24 28 8 12 3.29 0.02 13 5.35 -0.03 20 8 25 6.93 0.04 14 9.05 -0.13 15 8 19 7.49 0.04 15 10.56 -0.37 2 10 17 8.17 0.03 (a) = [K sq. deg. km/s] * = seem to form part of a larger complex

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71 TABLE 4.2 MOLECULAR CLOUDS IN THE CONNECTING ARM # <1> (°) (°) (km/s) V (km/s) r (’) : co (a) CT (a) 1 3.32 -0.52 260 27 15 8.46 0.05 2 3.98 -0.67 237 24 23 10.49 0.08 3 4.83 -0.87 206 29 27 22.3 0.1 4 5.45 -0.74 194 24 17 6.52 0.05 5 6.59 -0.68 158 22 34 23.7 0.1 6 7.77 -0.64 142 27 34 43.8 0.1 7 9.35 -0.91 120 16 40 11.3 0.1 (a) = [K sq. deg. km/s]

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72 TABLE 4.3 MOLECULAR CLOUDS IN FEATURE III b # <1> (°) (°) ( km/s V r ) ( km/s ) ( ' ) : co (a) CT (a) 1 4.33 0.01 89 10 12 1.99 0.03 2 6.76 0.07 118 12 12 1.06 0.03 3 7.56 0.05 122 10 13 2.30 0.03 4 8.32 0.00 126 15 13 3.83 0.04 5 9.20 -0.02 126 10 17 4.21 0.04 6 11.54 0.01 120 10 10 1.23 0.02 7 12.17 -0.02 118 17 13 2.46 0.03 8 12.80 0.01 101 8 10 0.78 0.02 (a) = [K sq. deg. km/s]

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73 TABLE 4.4 MOLECULAR CLOUDS IN THE 135 KM/S ARM # <1> (°) (°) ( km/s ) V r (km/s) (') X CO (a) cr (a) 1 -0.63 0.04 132 56 25 58.0 0.1 2 2.08 -0.15 133 41 24 36.1 0.1 3 2.45 0.29 159 15 15 3.67 0.04 4 5.34 0.26 182 7 22 2.60 0.04 5 5.89 0.33 190 10 15 1.10 0.03 6 7.92 0.29 183 15 19 3.99 0.05 7 8.94 0.36 181 17 22 4.06 0.06 (a) = [K sq. deg. km/s]

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74 TABLE 4.5 MOLECULAR CLOUDS IN FEATURE IX # <1> (°) (°) ( km/s ) V ( km/s ) r (') : co (a) (T (a) 1 5.92 0.01 137 12 13 3.15 0.04 2 6.60 0.03 141 5 9 0.43 0.01 3 6.70 -0.03 152 7 12 0.90 0.02 4 7.69 -0.06 153 7 12 1.21 0.02 5 10.65 -0.01 153 9 18 2.96 0.04 6 12.19 -0.08 147 10 13 0.96 0.03 7 12.73 -0.04 145 12 9 0.74 0.02 (a) = [K sg. deg. km/s]

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75 TABLE 4.6 MOLECULAR CLOUDS IN THE NORMA ARM # <1> (°) (°) V r ( km/s ) ( km/s ) ( ' r co ) (a) (T (a) 1 -11.30 0.165 -68 17 13 7.98 0.04 * 2 -10.82 0.130 -66 15 13 7.41 0.04 * 3 -9.95 0.082 -64 17 18 11.85 0.05 * 4 -8.49 0.252 -45 7 17 4.82 0.02 5 -8.00 0.-185 -53 12 15 5.40 0.03 6 -6.69 -0.040 -49 20 22 13.41 0.05 7 -4.36 0.690 -34 7 17 4.08 0.03 8 -1.99 -0.400 -24 7 17 5.47 0.03 * seem to form part of a larger complex (a) = [K sq. deg. km/s]

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76 TABLE 4.7 MOLECULAR CLOUDS IN THE SCUTUM-CRUX ARM # <1> (°) (°) (km/s) V (km/s) r (') I CO (a) CT (a) 1 -10.11 0.157 -39 12 21 10.54 0.04 2 -9.19 0.414 -36 15 13 4.88 0.03 3 -5.24 0.640 -21 7 17 10.94 0.04 (a) = [K sq. deg. km/s ]

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77 TABLE 4.8 MOLECULAR CLOUDS IN THE 70 KM/S ARM # <1> (°) (°) (km/s) V (km/s) r C) I CO (a) cr (a) 1 -11.58 0.02 4 12 10 1.29 0.02 2 -10.79 -0.01 15 10 10 1.55 0.02 3 -9.98 0.03 17 10 16 3.52 0.04 4 -8.33 0.05 42 10 10 0.66 0.03 5 6.44 -0.08 81 10 10 0.98 0.02 6 7.27 -0.03 91 10 17 2.99 0.03 7 10.99 -0.04 91 17 12 1.58 0.03 8 11.96 0.02 97 10 17 2.80 0.03 (a) = [K sq. deg. km/s]

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CHAPTER V THE KINEMATICS OF THE MOLECULAR GAS IN THE INNER GALAXY Introduction The kinematics of the atomic and molecular gas in the inner Galaxy are complex because of noncircular motions, large-scale asymmetries in longitude and velocity, and a spatial distribution inclined with respect to the Galactic plane. The situation is further complicated by the two-fold distance ambiguity inside the solar circle and by the pervasive effects of streaming motions of the gas. It is not surprising, therefore, that no interpretation or model proposed to date can successfully account for all the features observed. Most efforts towards understanding the kinematics of the inner Galaxy are based on H I data and, in general, fall into three broad categories, namely: 1) ejection models, 2) gravitational-field models, and 3) phenomenological models. Their characteristics and shortcomings are discussed below. In ejection models, the radial motions are attributed to gas in expansion expelled from the Galactic nucleus. An ejective origin of the 3 kpc and 135 km s -1 arms (Van der Kruit 1971), for example, would require the expulsion of about 10 8 M Q of gas from the nucleus (Oort 1977). Such 78

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79 large amounts of gas cannot be replenished by normal stellar evolution over the lifetime of the features (~ 10 6 yr) . Although this problem could be circumvented by circulation of the clouds through the center (Oort 1977), the essential remaining objection to ejection models is that several collimated events with various energies and occurring at different specific epochs have to be invoked to explain the observed features. Field models assume that the radial motions of the gas are the response to asymmetries in the gravitational field at the Galactic center. In order to explain the observed radial motions and the tilt of the gas distribution using the field hypothesis, the mass distribution at the center is required to be both nonaxi symmetric and tilted to the Galactic plane. While there is no evidence that the central star bulge--the main contributor to the mass density at the center--deviates from axial symmetry, a relatively small systematic deviation from symmetry is possible, and could explain some of the non-circular motions observed. For example, a 10% oval distortion of the bulge can account for the radial motions of the 3 kpc arm according to Yuan (1985). However, it is unlikely that the central star bulge has a significant tilt with respect to the plane, because the late-type stars of which it is composed have high random-velocities and are probably well-mixed. Moreover, inclined bulges are not observed in external edge-on

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80 galaxies, perhaps because the torques acting on an inclined bulge (Tremaine and Weinberg 1984) would make it unstable. Due to these constraints it is very unlikely that field models, except perhaps triaxial models (discused below), can account for the tilt of the gas layer and several other structures such as some jet-shaped features observed in H I, the molecular ring (Chapter II), and the wide-line clouds (Chapter III). Owing to the complexity of the gas motions in the Galactic center region, and the lack of knowledge of the dynamics involved, purely phenomenological models independent of dynamical considerations-that account for different features with a unifying assumption are useful to constrain dynamical interpretations of the inner-Galaxy gas. For example. Burton and Liszt (1978) were able to account for several observed H I features by assuming that most of the gas at the center is confined to a disk in expansion and rotation inclined 22° to the Galactic equator. Peters (1975) was able to account for the radial motions of the 3 kpc and 135 km s -*arms by assuming that the gas at the center moves in elliptical streamlines. Liszt and Burton (1980) have also accounted for several different H I features by assuming that the gas moves along elliptical streamlines such as those produced in a bar-like potential, but without reference to a particular mass distribution.

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81 The dynamical foundation lacking in phenomenological models could be obtained if one assumes that the stellar Galactic bulge is a rotating triaxial ellipsoid. Calculations of orbits in rotating triaxial potentials ( Schwarzchi Id 1979, Heiligman and Schwarzchild 1979) can produce closed elliptical orbits in retrograde motion and stable elliptical orbits inclined with respect to all the planes of symmetry. These orbits could explain, in principle, the presence of gas at "forbidden velocities" and the tilted distribution of the gas at the center (Lake and Norman 1982). So far, however, triaxial models have failed to reproduce the gas kinematics of the Galactic center. Heiligman (1982) used a triaxial potential, of the form shown to be self-consistent by Schwarzchild (1979), to calculate stable orbits intended to reproduce 13 C0 observations in the inner 400 pc of the Galaxy. The model, constrained to the plane, reproduces the general distribution of the gas in the 1-V plane, but not the 135 km s -1 arm or the molecular ring. Vietri (1986), developed realistic self-consistent models for the Galactic potential aimed at reproducing the inclined disk of Burton and Lizst (1978). He concludes that "...it is unlikely that the 22° tilt of the H I ring at r=1.5 kpc can be explained in terms of the tilted stable retrograde orbits in a rotating triaxial bulge," Vietri (1986, pg. 62).

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82 The failure of these different types of models to account for the kinematics of the gas near the Galactic center is probably a consequence of the use of oversimplified assumptions for the explanation of an intrinsically complex situation. In this state of affairs, it seems most useful to concentrate our efforts on accurately measuring observational parameters that can provide critical constraints for the different types of models discussed above. Owing to the high sensitivity, large spatial and velocity coverage and high sampling density of the present survey, the CO data collected in this survey are particularly suited for this purpose. The Molecular Disk Dynamical and phenomenological models of the Galaxy should account for the distribution and kinematics of the interstellar gas. Conversely, the characteristics of the molecular gas layer can provide important constraints on such models. In this section we discuss the measurements of the inclination and thickness of the molecular gas layer in the inner Galaxy obtained in this survey, and their relevance to models proposed for this region.

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83 The Tilt of the Molecular Disk Since the first H I surveys of the center region, observers have noted that the emission occurs preferentialy in opposing quadrants (1>0°, b<0° and 1<0°, b>0°), so that the gas distribution appears inclined to the Galactic equator. The magnitude and extent of the tilt of the gaseous disk are important quantities for models of the region. In particular, they are critical for the models proposed by Burton and Lizst (1978) and Lizst and Burton (1978, 1980), since their basic assumption is that the gas in the inner Galaxy (H I and CO) is contained in structures tilted to the Galactic plane. Although there is consensus on the direction of the tilt, its magnitude remains controversial. A tilt of about 8° was found by Kerr (1968) for the distribution of neutral hydrogen. The rotating and expanding disk proposed by Burton and Lizst (1978) is tilted by 22° with respect to h=0° and 78° to the plane of the sky. These values were found by fitting synthetic spectra to H I observations and from the inclination of high-velocity gas in moments maps. In another model, Liszt and Burton (1980) propose that the gas at the center is in pure rotation, confined to a barlike structure inclined 13.5° to the Galactic Equator, by 70° to the plane of the sky and at an orientation angle of 41.5°. These models are shown to account for H I features Ilia, IV, X, and XII (see Chapter II), and the authors claim

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84 that they also explain the distribution of molecular gas at the Galactic center. However, Heiligman (1982) observing CO within 2° of the center, found a tilt of 7° for the high-velocity molecular gas and no appreciable tilt for the high-intensity gas at the center. Sanders, Solomon, and Scoville (1984), also found a tilt of 7° from their 12 C0 observations for l>-4°, after subtracting the contribution of the gas in the foreground. The disagreement on the magnitude of the tilt of the gaseous disk at the Galactic center is due, in part, to different methods used to measure the angle of inclination. Discrepancies that remain after allowing for the different methods used probably reflect intrinsic differences in the distribution of the atomic and molecular gas at the Galactic center . In order to investigate how the methods used to measure the inclination affect the results, we have measured the tilt of the CO emission ridge to the Galactic plane in three different ways: a) from an 1-b map integrated over all observed velocities, i.e. from the distribution of CO emission in the plane of the sky, b) from 1-b maps integrated over selected velocity ranges that exclude the contribution of local emission, and c) from the 1-b distribution of the emission at extreme velocities. A problem encountered when attempting to measure the inclination of the molecular disk from 1-b maps is to define

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85 the structure or ridge whose inclination is to be measured. When the location and extent of the ridge is decided by eyeball estimates, considerable error occurs. In order to diminish this uncertainty, we adopted the following procedure, first we calculated the intensity-weighted mean latitude of the CO emission every 0.25° (2 beamwidths) in longitude, and then we fitted a least-squares intensityweighted straight line to the ridge of mean latitudes. We adopted this line as a first order representation of the CO emission in the plane of the sky, and used it to define the tilt of the molecular disk to the Galactic equator in the plane of the sky. Figure V-l shows the locus of the intensity-weighted mean latitudes, integrated over all observed velocities. The mean latitude ridge can be separated into 3 regions with different characteristics. In the central region (-1.5° <1< 1.8°), the ridge is virtually parallel to the Galactic plane, and the mean latitude slightly negative. At negative longitudes the ridge runs mostly above the plane reaching its maximum latitude of b=0.419° at 1=3.75°. At positive longitudes, conversely, the mean latitude ridge extends mostly below the Galactic equator reaching the minimum latitude of b=-0.387° at 1=4.75°. The large peak at 1=3.25° with mean latitude b=0.171° corresponds to the largest wideline cloud (see Chapter III).

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86 The straight line between 1=5° and l=-5° in Fig. V-l is the intensity-weigthed least-squares fit to the latitude ridge, in the region where the tilt is observed. The line is inclined to the Galactic equator by 2.3°, an angle smaller than all the values so far reported for the tilt of the gas distribution . at the center. This value, however, should be regarded as a lower limit because it was obtained integrating the emission over all observed velocities. Since this procedure overrepresents the contribution of the local gas, which is not likely to be inclined in a preferential way to the Galactic equator, the result is biased towards a smaller angle. The contribution of the inner-Galaxy gas can be separated from that of local material, in velocityintegrated 1-b maps, by a suitable choice of integration limits. Figures V-2 and V-3 present the distribution of the high-velocity molecular gas in the plane of the sky, respectively integrated from V=-319 km s -1 to V=-98 km s -1 and from V=98 km s ^ to V=319 km s ^ These limits were chosen because most local-gas emission occur outside these ranges, and because these limits allow direct comparison with H I observations (Burton and Liszt 1978). Our Figures V-2 and V-3 are the molecular counterparts of Burton and Liszt's H I moment maps (their Figs. 2 and 3). Comparison of Figs. V-2 and V-3 with Figs 2 and 3 of Burton and Lizst (1978) reveals that the inclination of the

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87 high-velocity molecular gas to the Galactic equator in the inner Galaxy is smaller than that of the high-velocity atomic gas. Although Burton and Lizst did not measure the inclination of H I emission to the Galactic plane directly from their figures 2 and 3, they argue that the inclination of the H I emission deduced from their H I moment maps support the tilt angle they propose for the H I disk (22° in Burton and Lizst 1978 and 13.5° in Lizst and Burton 1980). The variation of the intensity-weighted mean latitude of CO emission with Galactic longitude is displayed in Fig. V-4 for -319 < V <-98 km s -1 and in Fig. V-5 for 98< V <319 km s in the longitude ranges where significant emission takes place. To calculate the tilt of the molecular disk from the CO emission moment maps we have used the same technique described above, i.e. we have fitted an intensity-weighted least-squares straight line to the ridge of intensityweighted mean latitudes (Fig. V-6). The inclination of the lines are 8° for the negative velocities and 2.2° for the positive velocities. The smaller tilt found at positive velocities is due to the contribution of the wide-line cloud located at (l,b)= (3.2°, 0.3°), which produces a large fraction of the CO emission in this velocity range. It is possible that the wide-line cloud does not form part of the inclined molecular disk at the center. If the contribution of the wide-line cloud is subtracted, the

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88 inclination of the positive-velocity CO emission to the equator is 6°; similar to that found at negative velocities, and resulting in an average inclination (positive and negative velocities) of 7°. This value is in agreement with the angles measured by Kerr (1968), Heiligman (1982), and Sanders et al . (1984). The tilt angles calculated by these authors, therefore, have not taken into account the contribution of the wide-line cloud at (l,b) = (3.2°, 0.3°). Neither Kerr's nor Heiligman' s observations could have taken into account the contribution of the wide-line cloud. The lack of an H I counterpart to the wide-line cloud explains Kerr's result; while Heligman's survey, extending only to 1 = 2°, missed the wide-line cloud for over a degree. The Survey of Sanders et al. (1984) probably missed the wideline clouds due to undersampling. Their observations were spaced by 1° in 1, about twice the longitude extent of the largest wide-line cloud. In either case, however, the inclination of the CO emission to the Galactic equator calculated from moment maps is significantly smaller than the inclination of H I emission in analogous maps as reported by Burton and Liszt (1978) and Liszt and Burton (1980), and of CO emission as reported by Liszt and Burton (1978). The CO and H I moment maps also differ in the latitude extent of the respective emissions and their concentration to the central region. Most CO emission is confined between

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89 h=-l° and b=l° , while the H I emission extends to b=4° and b=-5°. This difference reflects the intrinsically smaller latitude scale-height of CO with respect to H I , and is not due to lack of latitude coverage of the CO emission. A low resolution CO survey of the center region, between b=-8° and b=8° (Nyman et al. 1987) confirms that essentially all CO emission from the center region falls within the latitude boundaries of the present survey. The emission from the region within a couple of degrees of the Galactic center is clearly distinguished from the rest of the inner Galaxy emission in the CO moment maps (Figs. V-2 and V-3). In the H I maps, on the other hand, the central region is blurred and confused with emission and absorption produced by surrounding gas. The inclinaton of the meanlatitude ridge is smaller in the central region (Fig. V-6 ) . We have calculated the inclination of the CO emission within 1° of the center, where the largest concentration of molecular mass is found. The inclination of the intensity-weighted least-squares fit is 2.5° for the positive velocities and 3.6° for the negative velocities. Comparison of the tilt of this central molecular disk with the inclinations found at larger longitudes suggests that the tilt of the CO emission to the Galactic equator increases with Galactic longitude. The velocity limits used in the moment maps still allow some contamination by emission which does not originate at

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90 the center; this affects the values found for the inclination of the molecular disk. In order to avoid any contamination, and obtain an upper limit for the tilt, only the highest velocities should be considered. The latitudes at which the highest velocities occur have been plotted as a function of Galactic longitude in Fig. V-7. The latitude distribution of the terminal velocities highlights the difference between the central region and the rest of the inner Galaxy. The high-velocity gas within ~ 1.75° (-250 pc) of the center is confined to a narrow band, parallel to the Galactic equator, at b=-0.125°. Outside this central region the latitude distribution of terminal velocities forms a band, extending from l=-6° to 1=5°, inclined to the Galactic equator by 9° for 1<0° and by 13° for 1>0°. Thus, the average inclination of the molecular disk, derived from the gas at terminal velocities, is approximately 11° ± 2°. It should be noted that the emission at higher inclination represents a small fraction of the emission (and thus of the molecular mass) tilted at smaller angles to the Galactic equator. This value is 4° larger than the inclination determined from the CO moment maps, and within 2.5° of the inclination of 13.5° calculated by Lizst and Burton (1980). However, the inclination deduced from the extreme velocities is significantly smaller than the 22° tilt found by Burton and Lizst (1978) and Lizst and Burton (1978). The extent of the CO tilted structure in Galactic longitude, is about 11°

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91 (r=800 pc). This value is significantly smaller than that of the H I disk proposed by Burton and Lizst (1978) (r=1.5 kpc), and that of the H I bar proposed by Lizst and Burton 1980) (1.9 kpc) . In summary, the picture that emerges from our study of the the tilt of the molecular gas in the inner Galaxy is one with four main components: 1) A sheet of high-density gas moving at high velocities within r ~250 pc, parallel to the Galactic equator at b~ -0.1°, with a thickness of less than 1 beamwidth (18.5 pc), and in circular motion around the center. 2) A central disk with a radius of about 250 pc and a half-thickness of 38±12 pc, inclined to the Galactic equator some 3° and to the plane of the sky by 84°, also in approximately circular motion. 3) A molecular gas layer extending some 5° at both sides of the center and inclined by 8° at negative velocities and 2.2° at positive velocities (or about 7° for positive and negative velocities, if the contribution of the largest wide-line cloud is disregarded), and 4) A disk or ring with a radius of about 5°, inclined some 11° to the Galactic equator and 87° to the plane of the sky, formed by gas at the terminal velocities. The division into components was done to facilitate the analysis and is not necessarily a physical one. Our observations are also compatible with a molecular disk whose inclination increases with Galactocentric radius, as suggested by Tohline et al . (1982).

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92 The total mass in the central sheet and disk (assuming purely gravitational motions) is about 4xl0 9 M Q , and the Q molecular mass less than 2.7x10 M Q (possibly an order of magnitude smaller, see Appendix A) . The inclination of the CO emission from the massive central component is about 3°, significantly smaller than values previously reported. The atomic gas, because of its larger scale-height and smaller concentration to the center, could appear at a \ larger tilt angle to the equator and extending to a larger radius. If this holds true, the H I bar inclined by 13.5° to the equator, by 70° to the plane of the sky and extending to a radius of ~2 kpc, proposed by Lizst and Burton (1980), could be an atomic gas counterpart to the CO terminal velocity gas inclined by approximately 11° to the Galactic equator. Structures further inclined to the Galactic plane, such as the 22° tilted disk of Burton and Lizst (1978), are most likely incompatible with our data. Further theoretical work on triaxial potentials would be useful to asses whether the inclination of the CO layer at the center is a consequence of the triaxiality of the bulge . The Thickness of the Molecular Disk As discussed in Chapter II most of the molecular gas of the inner Galaxy is found in a strip, approximately 2° wide.

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93 along the Galactic equator. In order to measure the thickness of this layer, perpendicular to the Galactic plane, profiles of velocity-integrated antenna temperature vs. Galactic latitude were produced every 0.25° (2 beamwidths) in Galactic longitude. The halfthickness of the molecular gas layer at each longitude was defined as the hal f -widthat-half -maximum (HWHM) of a gaussian curve fitted to the corresponding CO emission latitude profile. The average half-thickness of CO emission integrated over all observed velocities, north and south of the Galactic center, are very similar: 79 ± 20 pc in the 1 st quadrant and 81 ± 25 pc in the 4^ quadrant. Since the contribution of local gas was not removed in this calculation, the results are useful for comparison with tracers which do not have velocity information such as infrared and radiocontinuum. The halfthickness of the molecular disk at the Galactic center was calculated from latitude profiles after eliminating the contribution of the local gas, by integrating over appropriate velocity ranges. For longitudes between l=-5° and 1=0°, the latitude profiles were integrated between V=-320 km s" 1 and V=-100 km s" 1 ; while for longitude between 1=0° and 1=5°, the profiles were integrated between V=100 km s" 1 and V=320 km s" 1 . The average half-thickness of the molecular disk for negative longitudes was found to be 35 ±8 pc. While the

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94 corresponding value for positive longitudes was found to be 40±10 pc. The small difference north and south of the center cr) is not significant. Since the thickness of the gaseous disk is largely determined by the Galactic gravitational potential, the nearly identical values of the halfthickness of the molecular layer at both sides of the Galactic center are probably a consequence of axial symmetry of the Galactic bulge potential in a scale of several hundreds of parsecs. This is consistent with a bulge population formed by old, well-mixed stars. The half-thickness of the molecular disk at 1=0° is 37 pc. This value is 42% of the half-thickness of molecular material near the sun, 87 pc, found by Dame et al. (1987). If the disk of the Galaxy is modeled as a thin isothermal sheet, the FWHM of the gas layer (^ 1 / 2 ^ g ma Y b e approximated by (Van der Kruit 1983) 2 X / 2 1/2 (Z l/ 2 >g ~ 1 ' 7 < K > g [K G °(R)/Z 0 ] 7 (5-1) 2 1/2 where is the Z component of the gas velocity 9 dispersion, a(R) is the surface mass density, and Z Q is a scale parameter normal to the plane deduced from the light distribution of the disk (Z Q = 0.60±0.05 kpc). From (5.1) the surface density can be calculated if the thickness and the velocity dispersion of the gas are known.

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95 a(Z) 2.89 Z 71 G o (5.2) If the velocity dispersion of the gas at the center is similar to that observed near the sun, the difference in the thickness of the CO layer implies that the mass surface density at the center would be 5.5 times larger than that near the sun. Although we have not measured directly the velocity dispersion, we have indirect evidence that the velocity dispersion at the center is probably at least twice as large as that near the sun. We have measured the FWHM velocity width of the features at the terminal velocities. The widths are plotted vs. Galactic longitude in Fig. V-8. These profiles are widened by the instrumental profile, the doppler width of the clouds, and by the cloud to cloud velocity dispersion. At 1=1° and -1° the profiles are further widened by the Galactic differential rotation. The velocity widths at 1=11° to 1=13° are compatible with velocity dispersions in the outer Galaxy (Clemens 1985), while an envelope to the lowest values of velocity widths (to avoid the rotation broadening) gives velocity dispersions at least twice as large. Thus, a conservative lower limit for the mass surface density at the center is a , »22 a center sun In summary, the analysis of the thickness of the molecular disk favours models with an axisymmetric bulge

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96 potential, although it should be noted that a triaxial bulge can also produce the north-south symmetry observed in the CO disk thickness if the sun is properly aligned with an axis of the ellipsoid. The mass surface density at the center was found to be at least one order of magnitude larger than that near the sun. The Rotation Curve of the Inner Galaxy Knowledge of the rotation curve of our Galaxy, the variation of rotational velocity with Galactic radius, is important to constrain mass models of the Galaxy and to find kinematic distances to Galactic objects. A rotation curve of the inner Galaxy would be useful to see if any molecular gas in circular motion can be found, to compare the kinematics of the atomic and molecular gas at the Galactic center, and to compare the kinematics of the molecular gas north and south of the center. In this section we describe and analyze the construction and implications of a rotation curve of the inner 4 kpc of the Galaxy, derived from our CO observations. It should be noted, however, that the determination of the rotation curve of the inner Galaxy (r<4 kpc) is hindered by the presence of gas in non-circular motions, lack of axial symmetry, and gaseous structures inclined to the Galactic equator. Owing to these complications it has even

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97 been suggested that the determination of a reliable rotation curve in this region of the Galaxy may be intrinsically impossible (Knapp 1983). Even under these restrictions it is of interest to attempt the construction of the rotation curve of the inner Galaxy from CO observations. The derivation of the Galactic rotation curve is a fairly standard procedure (Schmidt 1965). In a Galaxy with perfectly circular rotation, the radial velocity, V , of an object in the plane at longitude "1" is (Mihalas and Binney 1981), V r = R o (W " W o> sind) (5-3) where R Q is the distance to the Galactic center, W is the angular velocity, , of the object, and W Q the angular velocity of the local standard of rest ( LSR) . The extreme velocity at each longitude corresponds to gas at the subcentral point along the line of sight. The subcentral point is located at Galactic radius R . min R sin(l) o ' ' (5.4) The extreme radial velocity (also called terminal or tangent velocity), is easily calculated by replacing (5.4) in (5.3), 0 sin(l) o ' Â’ 0(R) (5.5)

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98 The rotation curve is the variation of linear velocity with galactic radius 0(R), © ( R) = 0 O sin(l) + V t (5.6) This method requires knowledge of the constants R Q and 0 q and the determination of the terminal velocity at each longitude. For the constants we have adopted the values currently recommended by the IAU, R Q = 8.5 kpc and 0 q = 220 km s -1 ; to allow comparison with other works we have also calculated the rotation curve with the former IAU values (R Q = 10 kpc and © Q = 250 km s"-*-). The terminal velocities were determined directly from the spectra following the procedure outlined below. The survey data set (over 8000 spectra) was carefully examined to find emission features with extreme velocities (positive in the l s ^ quadrant and negative in the 4^ quadrant) every 0.5° (every 0.125° for |1| <2°), within 2° of the Galactic equator. The terminal velocities were calculated only from extreme velocity features with antenna temperature greater than 3a (0.36 K) and velocity span of at least 4 channels (5.2 km s ^). This procedure considers that the terminal velocity feature can arise from gas located outside the b=0° plane. The reasons for taking outof-the-plane profiles into account in the derivation of the Galactic rotation curve are discussed by Sinha (1978). The main requirement is that the Z-motion is effectively

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99 decoupled from motion in the plane of the Galaxy. If this condition is not met, out-of-plane terminal velocities may not correspond to circular rotation. It is interesting to note, however, that out-of-plane terminal velocities are compatible with a triaxial potential at the center (Vietri 1986) . There are several criteria to determine the point where the terminal velocity should be read in the velocity profile. We have adopted the criterion used by Kerr (1964) and Sinha (1978), in which the terminal velocity is measured at the point half-way up the steep slope of the extreme velocity feature. The advantages of this method over others used in this type of work are summarized by Sinha (1978). Before calculating the rotation curve, the terminal velocities were corrected to reflect the velocity of the cloud rather than its high-velocity wing. Clemens (1985), based on arguments relative to the line-width spectrum and cloud size spectrum, argues for a correction value of 3.0 km s . This correction value is very close to the velocity dispersion calculated by Alvarez et al. (1987) of 3.1 km s _1 , for observations of the 4 th quadrant obtained with the same telescope used in this survey. Therefore, we have adopted a nominal correction value of 3 km s -1 for our observations . The uncorrected terminal velocities for the I s *" quadrant (north) and the 4 th quadrant (south) are displayed

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100 as a function of Galactic longitude in Fig. V-9. Figure V-10 shows the rotation curve for R Q = 8.5 kpc and 0 Q = 220 km s _1 , while Fig. V-ll displays the same data for R Q = 10 kpc and 0 q = 250 km s _1 . The northern and southern rotation curves displayed as a function of sin 1 = R/R Q , for © Q = 220 km s _1 (Fig. V-12) show a similar general shape, although several differences are readily apparent. The northern curve rises linearly from V ~ 150 km s _1 , at the origin, to a broad peak between R/R 0 = 0.03 and 0.06, reaching V ~290 km s -1 , then decreases monotonically to V ~200 km s" 1 at R/R Q = 0.22. The southern rotation curve has a less steep, step-wise increase of the rotational velocity from the origin to R/R Q = 0.7. The southern maximum is narrower and slightly lower than the northern one. The rotation curve in the south decreases sharply to V ~110 km _1 at R/R Q ~0.1 and then increases to ~200 km s -1 at R/R Q = 0.2. A sharp peak at R/R Q =0.12 corresponds to an out-of-plane structure. The features that give rise to the terminal velocities can be easily identified in the integrated 1-V map of Chapter II (Fig. II-7). The terminal velocities in the north, between 1=13° and 1=6°, correspond to the 135 km s _1 expanding arm, and between 1=6° and 1=2.5° to the connecting arm. The terminal velocities within |1| <2° are produced by the nuclear disk. The southern peak is due to a structure symmetric to the connecting arm that extends down to 1=-

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101 4.5°. The rest of the southern curve originates in gas at the 3 kpc arm except at 1 ~ 7° where gas near b=0.5° produce a high velocity peak (see Figs. C-10 to C-13). To ease comparison between our rotation curve and other works, we have displayed our rotation curve superimposed to the H I rotation curve of Burton and Gordon (1978) (for R Q = 10 kpc and 0 Q = 250 km -1 ). The Burton and Gordon (B-G) rotation curve is widely used, specialy in the region of interest here (see e.g. Clemens 1986). Comparison between the B-G curve with the northern CO curve (Fig. V-13) reveals several differences. The maximum of the CO curve is approximately 30 km s -1 higher than that of B-G, it has more fine structure, and falls to lower velocities near R/R Q ~0.2. These differences can be easily understood. As seen in Fig. V7 the terminal velocities at r>250 pc occur outside the the b=0° plane, therefore, these terminal velocities were missed by in-plane surveys. The more abundant fine structure observed in our curves reflect the higher spatial resolution of our observations and the more clumpy distribution of CO compared to H I . The CO velocities near R/R Q -0.2 are lower than the B-G curve, but quite close to some of Clemens (1986) observational points which, in that region, lie systematically below the mean curve . An H I out-of-plane rotation curve (Sinha 1978) also shows higher peaks than the in-plane rotation curves, -290

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102 km s" 1 in the north and -270 km s" 1 in the south. The agreement of these figures with our results confirms that the higher peak velocities found in our work, as compared to in-plane surveys, are due to the inclusion of out-of-plane gas, and not to intrinsic differences in the kinematics of the tracers. Past the peaks, between R/R Q ~0.7 and R/R Q -0.2, the CO rotation curve is systematically lower than Sinha's and B-G's H I curves, specially in the south. The dif ferences may be the result of the higher concentration to the plane of CO emission compared to H I emission. The larger difference between the curves in the south (Fig. V14 ) may indicate that most of the molecular gas in this region is concentrated in the 3 kpc arm. The molecular gas within approximately 300 pc of the center, like the atomic gas (Rougoor and Oort 1960), is in nearly circular motion. If the motions are solely due to gravity, a total mass of approximately 4 x 10 9 M Q is contained in this region. It is interesting to note that a rotation curve very similar to the CO curve is produced by a triaxial bulge (Vietri 1986). Figure 8 in Vietri's paper shows a rotation curve that is very similar to our average (north and south) rotation curve (Fig. V-15). Vietri's rotation curve reaches a 290 km s ^ peak and a 190 km s ^ minimum, and corresponds to a triaxial bulge without rotation whose middle axis is aligned with the sun.

PAGE 113

103 Fig. V-l Locus of the intensity weighted mean latitude of CO emission in the plane of the sky, integrated over all observed velocities.

PAGE 114

p OJ cn CM 1 > >lO ^ rH II cn > 0 g o p p OH I OJ cn G 0 g rH ^ Qj OJ XS o P H P 0 Cn H Elj

PAGE 115

105 o r\j f\J — o — ou i i lanintn ouomao

PAGE 116

G 01 G PH cm ai 01 -P G g •H Ai ai II G > p Q) T 3 04 G g G ai -P r— I Gl G D) G 0) g 4-1 a: G G oo CT> O II U > >4 Ai g 01 O P (1) M-4 X! 4-1 » r— I mi 0 oi
PAGE 117

107 30nilib“l 31 iDHldD

PAGE 118

108 Fig.

PAGE 119

CfiLflCUC 109 Fig. V-5 Mean latitude ridge for velocities between V=100 km s” 1 and V=320 km s" 1 .

PAGE 120

110 Fig. V-6 Mean latitude ridge for positive and negative high velocities.

PAGE 121

Ill Fig. V-7 Distribution of terminal velocity features in the plane of the sky.

PAGE 122

VELDC I TY WIDTH (Kn/») 112 Fig. V-8 Velocity width of the terminal velocity features as a function of Galactic longitude.

PAGE 123

TERMINAL VELOCITr (Kn/ol Fig. V-9 Terminal velocities as a function of Galactic longitude. The squares and the triangles represent the I and IV Galactic quadrants, respectively.

PAGE 124

ROTRTION SPEED (Rn/sl 114 Fig. V-10 CO rotation curve for the inner Galaxy. The Galactic constants are 8.5 kpc and 220 km s" 1 . The squares and the triangles represent the I and IV Galactic quadrants respectively.

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ROTATION SPEEO lh./.l 115 Fig. V-ll CO rotation curve for the inner Galaxy. The Galactic constants are 10 kpc and 250 km s -1 . The squares and the triangles represent the I and IV Galactic quadrants respectively.

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ROT RT ION 116 Fig. -12 CO rotation curve for the inner Galaxy as a function of sin 1, for 220 km s" 1 . The squares and the triangles represent the I and IV Galactic quadrants respectively.

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117 Fig. V-13 CO rotation curve in the I quadrant compared to the H I curve of Burton and Gordon.

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ROTATION 6PEE0 lk n /„) 118 FigV-14 CO rotation curve in the IV quadrant compared to the H I curve of Burton and Gordon.

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119 Fig. V-15 Average CO rotation curve compared to the H I curve of Burton and Gordon.

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CHAPTER VI SUMMARY The first well-sampled, large-scale survey of ^CO (j=l_ > 0) emission from the inner 4 kpc of the Galaxy is presented, and used to study the distribution of molecular clouds and the kinematics of the molecular gas in the inner Galaxy. The survey samples a 4° wide strip along the Galactic equator from 1=-12° to 1=13°. The over 8000 spectra obtained with the Columbia University Southern millimeter telescope (La Serena, Chile) have a velocity resolution of 1.3 km s -1 , a rms sensitivity better than 0.12 K, and are spaced by approximately one beamwidth (8.8'). This is the first survey to encompass the complete latitude and velocity spans of the CO emission from the inner Galaxy. The survey is presented as a collection of 1-V maps at each observed latitude, 1-b maps integrated every 40 km s ^, and b-V maps at each observed longitude. Several new CO features were observed, and the molecular counterparts of the classic H I structures appear with unprecedented clarity owing to the dense sampling, high sensitivity, and extended latitude coverage of this survey. 120

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121 The two most striking features mapped in this survey exhibit remarkably large CO intensities and velocity widths. They are located at (1, b) = (3.2°, 0.3°), and (5.3°, -0.3°) , in a zone previously considered almost devoid of CO emission. Their respective CO luminosities exceed 8% and 3% of the CO luminosity of the inner 500 pc of the Galaxy, integrated over all observed latitudes and velocities. Their velocity widths reach 140 and 100 km s -1 (FWHM), respectively, the largest for molecular gas in the Galaxy beyond the nuclear region. We argue that these, and other wide-line features, are peculiar molecular clouds located in the vicinity of the Galactic center. Accordingly, their respective masses are 2.3 x 10^ and 6.4 x 10^ M Q , and their internal kinetic energies, 2.5 x 10 54 and 3.5 x 10 53 ergs. They appear to be gravitationally unbound, expanding in a time scale of approximately 1 x 10^ years. We argue that a plausible source for their large energies is a Seyfert-like event at the Galactic nucleus. The presence of such intense features in a zone previously considered devoid of CO emission, underlines the importance of well-sampled surveys in CO studies of the Galaxy. Most of the CO emission of the inner Galaxy is contained in layer approximately 2° wide in latitude, astride the Galactic plane, and extending over all observed

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122 longitudes. The CO emission is strongly concentrated in the inner 3.3° of the Galaxy. The molecular mass of this central region (A1 x Ab) = (490 x 190 pc), was found to be O 2.7 x 10 M q . The CO emission in this central region is distributed asymmetrically with respect to the Galactic center, with intensity-weighted mean coordinates: (1, b, V) = (0.4°, -0.05°, 23 km s”^). Over 60 molecular clouds were catalogued by their intensity-weighted mean coordinates. The catalogue also includes the velocity width, apparent radius and apparent CO luminosity of the clouds. The molecular clouds in the inner Galaxy trace the classic H I features remarkably well and with higher contrast than the H I observations. For example, portions of the Norma and Scutum-Crux arm were identified in CO a region wehere the the H I emission is too confused to make them out. The inclination and thickness of the molecular disk near the Galactic center were measured. Conflicting results on the tilt of the gaseous disk to the Galactic equator, cited in the literature, were found to arise partly from the different methods used to measure the angle of inclination and from differences in the distribution of H I and CO near the center. We found the molecular mass distribution to be tilted some 3° to the Galactic equator, within 1° of the center. The tilt of the larger scale component (R<5°) to the

PAGE 133

123 Galactic equator was found to be 8° for negative velocities and 2.2° for positive velocites. A consistent inclination of about 7° at both sides of the Galactic center is found if the contribution of the largest wide-line cloud is subtructed. The gas at terminal velocities was found to be further inclined to the Galactic equator, at an angle of 11° ± 2 °. The half-thickness of the molecular disk north and south of the Galactic center is very similar, 35±8 pc and 40±10 pc, favouring an axysymmetric or a preferentially oriented triaxial bulge model. The thickness and velocity dispersion of the molecular gas indicate that the surface mass density at the center is at least one order of magnitude larger than that near the sun. A rotation curve for the inner Galaxy was calculated from the CO terminal velocities within 2° of the plane, and compared to the H I rotation curve from Burton and Gordon (1978). The CO rotation curve is similar to that produced by a non-rotating triaxial ellipsoid (see Vietri 1986).

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APPENDIX A DETERMINATION OF MOLECULAR MASS IN CLOUDS CO Masses For the derivation of masses of molecular clouds from 1 2 their CO luminosities, we have followed the widely accepted practice of using the integrated 12 C0 line intensity, W CO J T < V > dv (A— 1 ) as a proportional tracer of the molecular gas in giant molecular clouds. Empirically W co has been found to be roughly proportional to the H 2 column density N(H 2 ) of molecular clouds in a wide range of physical conditions, despite the high optical depth of the 12 C0 line. This conclusion follows from the observation that line profiles of the optically thick 32 CO generally mimic the line profiles of the thin 13 C0 observed at the same positions. The theoretical basis for the proportionality between W C0 and N( h 2 ) is outlined by Scoville and Sanders (1987) in the following way: The 12 C0 luminosity, L co , of a molecular cloud can be defined as, L co = d2 J w co ^ = T co AV 71 r2 ( a ~ 2 ) 124

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125 where d is the distance to the cloud, Q is the solid angle subtended by the cloud, Tqq its peak brightness temperature, AV the line width, and R the cloud radius. For clouds in virial equilibrium. A V G M 1 2 ( A-3 ) so the CO luminosity, L co , can be written as. L CO 3 7i G 4 p ( A-4 ) where p is the mass density and M the mass of the molecular cloud. Since molecular clouds of different masses are found to have similar densities and temperatures, the CO luminosity, L co , is proportional to the cloud's mass. This proportionality holds for optically thick clouds with similar densities, because for such clouds the virial linewidth increases linearly with cloud radius, i.e. linearly with the column density of gas. Several different semiempirical techniques have been used to evaluate the proportionality constant between the CO emission and H 2 column density: ^CO emission measurements, visual extinction observations, calculations of virial masses, and gamma ray observations. The results of these independent techniques agree within a factor of two. Estimates of large-scale average molecular mass, hence.

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126 should be correct within a factor of two; mass calculations for individual clouds, however, may be less certain. For the determination of molecular masses in this survey, we have chosen the gamma-ray calibration value found by Bloemen et al. (1986) : N ( H 2 > 20 -211 — = 2.8 x 10^ mol cm K 1 km 1 s (A-5) CO because it is less likely than the values found with other techniques to depend on local conditions. Using this conversion factor, and a mean molecular weight per H 2 molecule of 2.76 x 10 gm (Allen 1973), the masses I O obtained from ^CO luminosities can the be expressed as, M co = i9 x 10 3 I co d 2 (A-6) where Mqq is in solar masses, d is the cloud's distance in kpc, and I co is the apparent CO luminosity of the cloud, i . e . , ^0 = n t dv ^ ( a— 7 ) with AV in km s ^ and Q in square degrees. LTE Masses The masses of molecular clouds can also be estimated 12 13 from CO and CO observations assuming local thermodynamic equilibrium (LTE), optically thick 12 C0 lines, optically

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127 13 thin CO lines, and adopting a constant ratio between 33 CO and H 2 column densities. Under the assumption of LTE, the thermal motions of the H 2 molecules and the 12 C0 and 13 C0 rotational levels are in equilibrium, so kinetic and excitation temperatures are identical (T x = T k ) . For the optically thick 12 C0 line, the line antenna temperature, T L , is related to the excitation temperature T x by the expression (see Blitz 1978) T x = 5.53 (In (1 + ( 5.53 T r + 0.819) ±2 )) -1 ( A-8 ) and the optical depth of the thin 13 C0 line is given by: t 13 = _ln(1 ' ( 5 tH' ) (( ex P(-^^-)-D“ 1 -0.164)" 1 ) ( A-9 ) From these quantities, the 13 C0 column density, N( 13 C0), is found: N( 13 C0) = 2.42 x 10 14 < T 13 AV V ( 1-exp ( ~ 5 ' . 29 ) ( A-10 ) where AV is measured in km s” 3 . 13 The CO column densities are converted to H 2 column densities using the proportionality factor found by Dickman (1978), N( 13 C0) N(H 2 ) 2 x 10“ 6 (A11 )

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128 using a molecular weight per H 2 molecule of 2.76 x 10“ 24 gm (Allen 1973), the LTE mass of a cloud can be written as M LTE = 2 19 x 10 A N(H 2 ) (A-12) where M LTE is expressed solar masses, A is the area of the cloud in pc , and N(H 2 ) is the H 2 column density in cm'^. Masses calculated with the LTE method are lower limits because the 32 CO cloud limits are larger than the 42 CO limits. The accuracy of this method is critically dependent on the consistency and value of the ratio between 13 C0 and H 2 column densities. Virial Masses An independent value for a cloud's mass can be obtained assuming that the cloud is in virial equilibrium supported by internal motions. The virial theorem for a uniform, spherical cloud of mass M, radius R, and 3-dimensional velocity dispersion a, can be written as 2 _ (3 G M a = 2 M ) 5 R ( A13 ) so the cloud's virial mass can be expressed as M vir = 210 R AV 2 ( A-14)

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129 where is in solar masses, R is in pc, and AV is the FWHM velocity width of the cloud in km s -1 . The virial mass is an upper limit to the cloud's mass because the effects of magnetic fields and density gradients have been ignored. Molecular Masses near the Galactic Center It is important to note that conventional estimates of H 2 mass in the Galactic center using CO observations may be in error by as much as one order of magnitude. As Oort (1977) first pointed out, when the large CO luminosities observed towards the center are combined with the abundance ratio N(H 2 )/N( 12 CO) (e.g. Gordon and Burton 1976) and with the gas-to-dust ratio observed in the neighborhood of the sun (Savage and Mathis 1979), visual extinctions of 150 to 200 magnitudes in front of the nucleus are derived (Oort 1977, Blitz and Shu 1980). Photometric observations at various wavelengths give visual extinctions of about 27 to 30 magnitudes in front of the central source (Willner 1976, Becklin and Neugebauer 1969), about half of which is produced in the disk (Oort 1977). Therefore, the observed visual extinction of the Galactic center is an order of magnitude smaller than that derived from CO luminosities and standard conversion factors. Observations of gamma-rays from the Galactic center indicate a flux value an order of magnitude smaller than what should be expected if the H 2 masses derived from CO

PAGE 140

130 data were correct (Blitz et al. 1985). Thus, these observations also support a smaller mass at the center than that derived from CO observations and standard conversion factors . Plausible reasons for the discrepancy between observed and calculated visual extinctions are that the conventional abundance ratio and/or the gas-to-dust ratio do not apply at the Galactic center. Since this is not yet clear, nor have more pertinent values for the conversion factors been derived, we will use the standard conversion factors with the caveat that the real masses may be an order of magnitude smaller .

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APPENDIX B LONGITUDE -LATITUDE MAPS This appendix presents the data set in 1-b maps integrated every 40 km s -1 , over a velocity range of 600 km s" 1 . These "velocity slices" are useful to identify molecular clouds and kinematic structures. The contours are spaced by 20 K km s \ and the lowest contour is set at 5 K km s“^. 131

PAGE 142

132 aoruutn jiiDdieo

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133 3onuitn oiiDbidD

PAGE 144

134 o Z) ID 2 O l_> cr cr o ro ! Cn •H aanmtn ouotnao

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135 3onmtn DiiDtndo

PAGE 146

136 aoruiim onotnoo

PAGE 147

137 30D1 1 101 DUDblbO

PAGE 148

138 aanuitn diiduiuo

PAGE 149

139 aanuitn 31130100

PAGE 150

140 30(11 1 lti“l DUDblbO

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141 e ru » o o o O r\J — o — . ru i i aanintn onaoioo

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142 LU CD 3 LD Z o CJ (_) cr _j a o I CQ (T> •H & 30D1I1U1 31 13dlt)D

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143 CM 00 — • O — 00 I I CN I — I I P 3 tP •H 30011101 GUDtildO

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144 3oni.ij.ai Diiodiao

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145 i — i i ra • Cn •H u< 30nillbl DIlDdldO

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146 LO i — i i 03 CP •H fa aanmtn diidhiod

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APPENDIX C LONGITUDE VELOCITY-DIAGRAMS This appendix presents the data set in 1-V diagrams at each observed latitude. This representation of the data is particularly useful to identify kinematic structures and follow their development with Galactic latitude. The contours are separated by 1 K, and the lowest contour is set at 0.5 K. 147

PAGE 158

LV DIRGRfiM B = 2.500 148 aaruiONOi aiioyibo -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RROIRL VELOCITY (KM S ')

PAGE 159

LV DIRGRRM B = 2.250 149 t 1 I I — — I I 3QniI0N01 DUDtntiQ -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VELOCITY (KM S')

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LV DIRGRRM B = 2.000 150 I I I I — — I I BOnilONOl 31 IDdldO -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RRDIHL VELOCITY 1KM S' 1 )

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LV DIAGRAM B = 1.750 151 (_) o Lu > Q cr cc cc co i i I U • tn •H fa 3QniI0N01 DIlDbldO

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LV DIAGRAM B = 1.500 152 ‘ I 1 l 1 l 1 I 1 i 1 I 1 I ' I ' I 1 I 1 I '' I 1 I ' I 1 I 1 I 1 I 1 I 1 I 1 I CO LkJ cc & 'O & .J8 o


PAGE 163

LV DIAGRAM B = 1 .250 153 I I I I ~ — I I 3001 1 0ND1 OUDbltiO -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VLL0C1 TY (KM S' 1 )

PAGE 164

LV DIAGRAM B = I .000 154 3onnoNcn aiiDbibo -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR ROD I OL VELOC 111 (KM S' 1 )

PAGE 165

LV D I RGRflM B = 0.880 155 aaruiONtn oiiDtiidO -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 L,SR RADIOL VELOCITT IKM S')

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LV DIAGRAM B = 0.750 156 I 1 i 3 —i O cn .'A A , 1 o e __°©o©oeo ojococD3*c\joc\j=r CO I e o CD O I rsj i laniioNDi ouibibo “300 -250 -200 -150 -100 -50 0 50 100 150 200 250 LSR RADIAL VELOCITY (KM S' 1 )

PAGE 167

LV DIAGRAM B = 0.630 157 jOnilONGI 31 lDdldO -300 -260 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VELOCITY (KM S')

PAGE 168

LV DIAGRAM B = 0.500 158 i i i i — — i i 3QniI0NCn DIlDblbG -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 L5R RflDIRL VEL0C1 TT (KM S' 1 )

PAGE 169

LV D I RGRflM B = 0.380 159
PAGE 170

LV DIAGRAM B = 0.250 160 30m lONcn oiiDbibo -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VELOCITY (KM S')

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LV DIAGRAM 161 3aniIDN01 DIlDblbD -300 -250 200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VELOCITY (KM S ')

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LV DIAGRAM B = 0.000 162 aaruiONOi diidbibo

PAGE 173

163

PAGE 174

LV DIAGRAM B = -0.250 164 3arUI0NQ1 jIlDtHdO -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIOL VELOCITY (KM S' 1 )

PAGE 175

LV DIAGRAM B = -0.380 165 CO I — I I u Cn •H aaniioNcn aiiDbibo

PAGE 176

LV DIAGRAM B = -0.500 166 ianiioNoi oiiobibo -300 250 200 -150 -100 -50 0 50 100 150 200 250 300 LSR RAO I AL VEL0C1 H (KM S' 1 )

PAGE 177

LV DIAGRAM B = -0.G30 167 30ni I ONDl 31130300 -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 LSR RAO I AL VELOC I TT (KM S ')

PAGE 178

LV DIAGRAM B = -0.750 168 30niI0\01 OIlDblbG -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VEL0C1TT IKM S')

PAGE 179

LV DIAGRAM B = -0.880 169 aannoNoi ouobibo -300 -280 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VELOCITY IKM S"')

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LV 0 1 fiGRRM B = -1.000 170 o o m o LD OJ o o Pu o LO o o XL o — LO LJ o > — J cr o o o in cc ro CN I U •H CM lamiDNOi oiiotnuo

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LV DIAGRAM B = -1.250 171 I i aaniioNci oi lObioo -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL VELOCITY (KM S' 1 )

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LV DIAGRAM B = -1.500 172 lOnilONOI DUDblbO -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RAOIAL VELOCITY (KM S' 1 )

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LV DIflCRRM B = -1.750 173 In 31 id LJ O o qc co CO cn i i VO CN I U b'l •H 3GniI0N01 DIlDbldO

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LV DIAGRAM B = -2.000 174 cn -XL (_> o Q C E ac cc CO I I (N l U Cn •H &4 3Qfli I GN01 DllOblbO

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LV DIAGRAM B = -2.250 175 j.' 1 ' n ' i ' i ' i r r 1 i 1 i 1 i '' i 1 i 1 i 1 i 1 i i 1 i 1 i 1 i ' i 1 i ' i ' i i aoruioNcn diidbioo 300 -250 200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIAL Vt'LOCITY (KM SÂ’ 1 )

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LV DIAGRAM B = -2.500 176 cD?o , 1 . 1 , 1 . 1 . 1 . 1.11 ° O o o o O o O CO CD OJ O f\J I CD I 3aniIQNGl DliOtnbO -300 -250 -200 -150 -100 -50 0 50 100 150 200 250 300 LSR RADIOL VELOCITY IKM S' 1 )

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APPENDIX D LATITUDE -VELOCITY DIAGRAMS The latitude structure of the CO emission is presented in these b-V diagrams plotted at each observed longitude. The contours are separated by 1 K, and the lowest contour was set at 0.5 K. The velocity and latitude limits of the observations are indicated by a dotted line. 177

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GOLRCT I C LRTITUDE 178 L = 1 3! 0 U 12! 875 -300 -200 -100 0 100 200 300 L = \z‘. 75 Fig. D-l

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GALACTIC LATITUDE 179 L = 12! 5 L 12! 375 L = 12! 25 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-2

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GAL ACT I C LATITUDE 180 -300 -200 -100 0 100 200 300 L 1 1 ! 625 LSR RADIAL VELOCITT (KM S' 1 ) Fig. D-3

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GflLflCTIu LRTITUDE 181 -300 -200 -100 0 100 200 300 L 1 l! 125 LSR RADIAL VELOCITY (KM S") Fia. D-4

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GflLRCT I C LRTITUDE 182 L 1 o! 875 L 1 0! 625 LSR RflDIRL VELOCITY (KM S' Fig. D-5

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GAL AC T I C LATITUDE 183 L = 1 o! 5 -300 -200 -100 0 100 200 300 l 10! 125 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-6

PAGE 194

GALACTIC LATITUDE 184 L = 1 o! 0 L = 9?75 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-7

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GALACTIC LATITUDE 185 L * 9! 5 L 9! 375 L = 9? 25 l * 9! 125 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-3

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GALACTIC LATITUDE 186 L = 9!0 L 8! 875 L 8.625 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-9

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GALRC T 1 1 _HT I TUDE 187 l = e!s L * 8! 375 L = 8! 25 L 8! 125 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-10

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GHLRC TIC LRTITUDE 188 L * b! 0 L 7?875 L * 7! 75 L 7! 625 LSR RflDIRL VELOCITY (KM S' 1 ) Fig. D-ll

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GflLflC'i .0 LRTITUDE 189 L = 7! 5 2 ° 1 ° 0 ° 2 ° -300 -200 -100 0 100 200 300 2 ° 1 ° 0 ° 1 ° 2 ° -300 -200 -100 0 ICO 200 300 L 7! 125 -300 -200 -100 " 0 100 200 300 LSR RPDIflL VELOCITY (KM S* 1 ) L = V.25 L * 7. 375 Fig. D-12

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GALACTIC LATITUDE 190 L = t!o L 6! 875 L = 6! 75 L 6! 625 -300 -200 -100 0 100 200 300 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-13

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GALACTIC LATITUDE 191 L = 6^5 l 6! 375 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-14

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GALACTIC LATITUDE 192 l = e!o L 5! 675 L = 5? 75 L 5! 625 -300 -200 -100 0 100 200 300 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-15

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GALflC T I C LATITUDE 193 L = 5? 5 L 5! 375 L = 5! 23 L * 5! 125 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-16

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GALACTIC LATITUDE 194 L = 5^0 l 4? 875 L = A! 75 L 4! 625 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-17

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GRLflC : LATITUDE 195 L = U!5 L • 4? 375 L = U?25 L * m! 125 LSR RADIAL VELOCITY (KM S"') Fig. D-18

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GALACTIC LOT I TUDE 196 i * u!o L • b! 875 L • 3! 625 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-19

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GPL AC T I C LATITUDE 197 L = 3? 5 L = 3?2S LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-20

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GPLfiC TIC LRTITUDE 198 L * b!o L = 2! 75 L 2! 62S LSR RRDIRL VELOCITT (KM S") Fig. D-21

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GALACTIC LATITUDE 199 L = 2? 5 L 2°. 375 L = 2° 25 L 2'. 125 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-22

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GALACTIC LATITUOE 200 L = 2°0 LSR RADIAL VELOCITT (KM S' 1 ) Fig. D-23

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GALAC1 IC LATITUDE 201 L = l!s -300 -200 L l! 375 100 100 200 L * 1.25 L IT 125 300 LSR RADIAL VELOCITY (KM SÂ’ Fig. D-24

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GflLACTIi. LATITUDE 202 -300 -200 -100 0 100 200 300 L o! 875 L * 0?75 L O! 625 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-25

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GAL ACT I C LATITUDE 203 L • o!s LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-2S

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GALRCT i _ LATITUOE 204 L « o!o L -o! 125 l = -0! 25 L • -O! 375 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D27

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GALACTIC LATITUDE 205 l -o! S25 L = -O! 75 L -Ol 875 -300 -200 -100 0 100 200 300 LSR RADIAL VELOCITT (KM S' 1 ) Fig. D-28

PAGE 216

GALACTIC LATITUDE 206 L * -i!o L -l! 125 L 1 ! 375 LSR RADIAL VELOCITY (KM S") Fig. D-29

PAGE 217

GALACTIC LATITUDE 207 L = -l! 5 L — 1 ! 625 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-30

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GALACTIC LATITUDE 208 l • -2! 125 LSR RADIAL VELOCITY (KM S") Fig. D-31

PAGE 219

GALACTIC LATITUDE 209 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-32

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GALACTIC LATITUDE 210 L = -3!0 L -3! 125 L = -3! 25 L -3! 375 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-3 3

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GALACTIC LATITUDE 211 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-34

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GAL AC T I C LATITUDE 212 L • -I*: 125 L = -u! 25 L -n! 375 LSR RADIAL VELOCITY (KM S") Fig. D-35

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GALACTIC LATITUDE 213 -200 -100 0 100 200 300 L -w! G25 -300 -200 -100 0 100 200 300 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-36

PAGE 224

GALACTIC LATITUDE 214 L = -5!0 L -5! 125 L = -5! 25 L • -5! 375 LSA RADIAL VELOCITY (KM 5") Fig. D-37

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GALACT I 1LATITUDE 215 L = -5*. 5 L -5! 625 L = -5! 75 L • -5! 875 LSR RADIAL VELOCITY (KM S'*) Fig. D-38

PAGE 226

GALACTIC LATITUDE 216 L * -6!0 L * -6! 25 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-39

PAGE 227

GALACTIC LATITUDE 217 L = -6! 5 L -6! 625 L * -6! 75 L -6! 875 LSR RADIAL VELOCITY (KM S") Fig. D-4 0

PAGE 228

GRLRCT I C LATITUDE 218 L = -7! 0 L -7! 125 l = -?!2S L -?! 375 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-4 1

PAGE 229

GRLRC T I C LATITUDE 219 LSR RADIAL VELOCITY (KM S") Fig. D-42

PAGE 230

GALACTIC LATITUDE 220 LSR RADIAL VELOCITY (KM 5'') Fig. D-43

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GALAC'i , LRTITUDE 221 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-44

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GALACTIC LATITUDE 222 LSR RADIAL VELOCITY (KM S' Fig. D-43

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GALACTIC LATITUDE 223 L * -9! 5 LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-4S

PAGE 234

GRLACT I C LATITUDE 224 L * -io!o L -10? 125 LSR RflDIfll VELOCITY (KM S' 1 ) Fig. D-47

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GALACTIC LATITUDE 225 L = -10?S L 1 0.625 L -101875 LSR RADIAL VELOCITY (KM S' 1 : Fig. D-43

PAGE 236

GALACTIC LATITUDE 226 i = -n!o LSR RADIAL VELOCITY (KM S' 1 ) Fig. D-49

PAGE 237

GRLRC1 LATITUDE 227 LSR RRDIRL VELOCITY (KM S' 1 ) Fig. D-50

PAGE 238

228 L = — 1 2 ? 0 Fig. D-51

PAGE 239

REFERENCES Allen, C. W., 1973, "Astrophysical Quantities," ( London : Athlone ) Allen, D. A., Hyland, A. R. , Jones, T. J. , 1983, M.N.R.A.S., 204, 1145. Altenhof f , W. J. , Downes, D. , Pauls, T. , and Schraml, J. , 1978, Astr. Ap. Suppl., 35, 23. Alvarez, H. , May, J., and Bronfman, L., 1987, in preparation. Audouze, J., Lequeux, J., Masnou, J.-L., and Puget, J.-L., 1977 Astron. Astrophys . , 80, 276. Bania, T. M. , 1977, Ap. J., 216, 381. Bania, T. M. , 1980, Ap. J. , 242, 95. Bania, T. M. , 1986, Ap. J. , 308, 868. Bania T. M. , Stark, A. A., and Heiligman, G. , 1986, Ap. J. , 307, 350. Becklin, E. E., and Neugebauer, G. , 1969, Ap. J. (Letters), 157, L31 . Bieging, J., Downes, D., Wilson, T. L. , Martin, A. H. M. , and Gusten, R. , 1980, Astron. Astrophys. Suppl., 42, 163. Bitran, M. , Alvarez, H. , Cohen, R. S., and Thaddeus, P., 1985 Bull Am. Astron. Soc . , 17, 607. Blitz, L. , 1978, Ph.D. dissertation, Columbia University, published as NASA Tech. Memo. No. 79108, 1978. Blitz, M. , L., Bloemen, J. 1985, Astr. Ap. B. G. M., , 143, 267 Hermsen, W. , and Bania, T Blitz, L . , and Shu , F . H., 1980, Ap. J. , 238, 148. Bloemen, J. B. G. M. , Strong, A. W. , Blitz, L., Cohen, R. S., Dame, T. M. , Grabelsky, D. A., Hermson, W. , Lebrun, F., and Thaddeus, 1986, Astron. Astrophys., 154, 25. 229

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BIOGRAPHICAL SKETCH Mauricio Ernesto Bitra'n Carreno, the oldest son of Raul Bitran (Z.L.) and Cora Carrefio de Bitran, was born the 10^ of July 1954 in La Serena, Chile. He is married to Gloria Rachamin and they have one son, Jonathan Abraham Bitran. After Mauricio finished elementary school in the "Escuela Superior de Hombres # 1" in La Serena, the family-parents, sisters Jazmin and Marcela, and brothers Leonardo and Cl audio-moved to Santiago, the capital of Chile. There he completed his high school studies at the "Liceo Manuel de Salas," after which he entered the "Facultad de Ciencias Fisicas y Matematicas" of the University of Chile, where he obtained the degrees of "Bachiller en Ciencias" (Physics) and "Magister en Ciencias, con distincion maxima" (Astronomy) with the thesis "La radiacion del fondo Galactico en 45 MHz." In 1981, he started graduate studies at the Department of Astronomy of the University of Florida, obtaining a Master of Science degree in 1984. A Ph.D. degree is expected to be awarded in December 1987. 235

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I certify that I have read this study and that in my opinion it conforms to acceptable standards of scholarly presentation and is fully adequate, in scope and quality, as a dissertation for the degree of Doctor of Philosophy. c, . Stephen T. Gottesman, Chairman Professor of Astronomy I certify that I have read this study and that in my opinion it conforms to acceptable standards of scholarly presentation and is fully adequate, in scope and quality, as a dissertation for the degree .of Doctor of Philosophy. Philos (, !/{ i o~ d Lc'-i Patrick Thaddeus, Cochairman Professor of Physics I certify that I have read this study and that in my opinion it conforms to acceptable standards of scholarly presentation and is fully adequate, in scope and quality, as a dissertation for the degree of Doctor of Philosophy. Thomas D. Carr Professor of Astronomy I certify that I have read this study and that in my opinion it conforms to acceptable standards of scholarly presentation and is fully adequate, in scope and quality, as a dissertation for the degree of Doctor of Philosophy. u . I r— Jdmes Hunter, v Jr. Professor of Astronomy

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I certify that I have read this study and that in my opinion it conforms to acceptable standards of scholarly presentation and is fully adequate, in scope and quality, as a dissertation for the degree of Doctor of Philosophy. of Astronomy HaVwood Smith, tar Associate Profess I certify that I have read this study and that in my opinion it conforms to acceptable standards of scholarly presentation and is fully adequate, in scope and quality, as a dissertation for the degree of Doctor of Philosophy. JU.~. Lltffc-.. William Weltner Professor of Chemistry This dissertation was submitted to the Graduate Faculty of the Department of Astronomy in the College of Liberal Arts and Sciences and to the Graduate School and was accepted as partial fulfillment of the requirements for the degree of Doctor of , Philosophy . December 1987 Dean, Graduate School