A Dissertation entitled
The Transition From Diffuse to Dense Molecular Clouds
by Johnathan S. Rice
Submitted to the Graduate Faculty as partial fulfillment of the requirements for the Doctor of Philosophy Degree in Physics
Dr. Steven R. Federman, Committee Chair
Dr. B-G Andersson, Committee Member
Dr. Song Cheng, Committee Member
Dr. Tom Megeath, Committee Member
Dr. John-David Smith, Committee Member
Dr. Amanda C. Bryant-Friedrich, Dean College of Graduate Studies
The University of Toledo Dec 2018 Copyright 2018, Johnathan S. Rice
This document is copyrighted material. Under copyright law, no parts of this document may be reproduced without the expressed permission of the author. An Abstract of The Transition From Diffuse to Dense Molecular Clouds by Johnathan S. Rice
Submitted to the Graduate Faculty as partial fulfillment of the requirements for the Doctor of Philosophy Degree in Physics The University of Toledo Dec 2018
The atomic to molecular transitions occurring in the diffuse interstellar gas sur- rounding molecular clouds are affected by the local physical conditions (density and temperature) and the radiation field penetrating the material. Our observations at visible wavelengths of CN, CH, CH+, Ca I, Ca II, and K I absorption from McDonald
Observatory and the European Southern Observatory are useful tracers of this gas and provide the velocity component structure needed for analyzing lower resolution, archival ultraviolet observations of CO and H2 absorption from the Far Ultraviolet Spectroscopic Explorer. Additional archival data from ultraviolet measurements of
CO and C I absorption were acquired with the Hubble Space Telescope to complement other data sets.
A relationship between the column densities of CO and H2 has been known for some time. The break in the slope of the power law relation between CO and H2 column densities corresponds to a change in the photochemistry. Our CO and H2 column density measurements strengthen the results from previous surveys over the ranges corresponding to the transition between diffuse molecular gas to translucent gas. These trends are used in conjunction with those seen in other species like CH and H2 to infer information about the physical conditions in the gas. The derivation of densities can provide additional insight into the correct interpre- tation of trends found in our data and other surveys. We further explore the changing iii environment between diffuse and dense gas by using the column densities and excita-
tion temperatures from CO and H2 to determine the gas density. The resulting gas densities from this method are compared to densities inferred from other methods such as CN chemistry or C2 and C I excitation. The total proton gas densities (nH
= n(H I)+2n(H2)) allow us to find trends in the combined set of tracers.
• Densities from C I and C2 excitation are typically smaller than the densities
−3 from other tracers. (nH < 100 cm )
• Larger densities are derived from 12CO(J = 1-0) excitation and are similar to
+ −3 −3 those from C excitation and CN chemistry. (100 cm < nH < 300 cm )
• The largest densities are derived from 12CO(J = 2-1) excitation and are typically
−3 larger than those from CN chemistry. (nH > 300 cm )
Additional kinematic associations of the different tracers with H I,C+, 12CO, 13CO, and C18O emission from the GOT C+ survey extend the combined picture from other tracers and strengthen the characterization of diffuse interstellar gas, including
CO dark gas, which is not associated with emission from H I at 21 cm or from CO at
2.6 mm. From the kinematic associations, we find
• Diffuse Atomic Clouds have components seen in Ca II, Ca I, H I, and sometimes C+;
−3 (nH < 100 cm )
• Diffuse Molecular Clouds (CO Dark Gas) have components seen in C+, K I, CH+,
12 13 −3 −3 CH, or CN but without CO or CO emission; (100 cm ≤ nH ≤ 300 cm )
• Dense Molecular Cloud Envelopes have components seen in 12CO and 13CO emission
(with the distinction that clouds with 13CO are denser) and any of the additional
−3 −3 species seen in emission or absorption. (300 cm ≤ nH ≤ 1000 cm )
Groupings of sight lines, such as those toward h and χ Persei or Chamaeleon provide a chance for further characterization of the environment. The Chamaeleon Region in particu- lar helps illuminate CO dark gas. This complements other dark gas searches with OH, C+, iv and H I Narrow Self Absorption (HINSA) features. The results described here are the latest effort seeking details of the correspondences revealed by the earlier studies and attempt to weave the results from different methods and tracers into a more complete picture.
v This manuscript is dedicated to my family: my dad, Scott Eugene Rice, who has always been my hero, my mom, Robin Diane Rice, who has always guided me through the good and bad times, my brother, Martin Alfred Rice, who left this world too soon, and my wife, Nicole Elizabeth Cappelletty, within who I see my future. Acknowledgments
The completion of this Dissertation would not have been possible without the patience and understanding of my advisor, Dr. Steven Federman. His input and patience with the numerous revisions of this Dissertation, and all the work leading up to it, is something for which I will always be thankful. I would also like to thank him for giving me encouragement when I needed it, all the help and understanding during the hard times, and all the advice he has had to repeatedly give me over the years.
I have also had the help of many collaborators over the years and this work would not be what it is today without their help: Dr. Paul Goldsmith, Dr. William
Langer, Dr. Jorge Pineda, Dr. Nicolas Flagey, Dr. David Lambert, Dr. Adam Ritchey,
Justin Brown, Dhruv Sengar, Malinda Bender, and Johnathan Winckowski. I also would like to thank the other members of my committee, Dr. Song Cheng, Dr. Tom
Megeath, Dr. JD Smith, and Dr. B-G Andersson, for their support and feedback. I am grateful to Yaron Sheffer for allowing me the use of his program ISMOD. This research made use of the SIMBAD database operated at CDS, France. Some of this work was performed at the Jet Propulsion Laboratory, California Institute of
Technology, under contract with NASA. This work was supported in part by NASA
Grant NNX10AD80G. Additional support, for the work done by Dhruv Sengar and
Malinda Bender in the REU program, came from Grants Phy-1004649 and Phy-
1262810, respectively.
I would like to thank my parents for not only their help and support but for inspiring me to pursue my dreams and become the person I am today. Lastly, I
vii would like to thank my wife, Nicole, for all the support, understanding, and endless patience you have given me since the day I met you. You have changed my life more than I can possibly express.
viii Contents
Abstract iii
Acknowledgments vii
Contents ix
List of Tables xiii
List of Figures xviii
1 Introduction 1
1.1 ISM Overview ...... 1
1.2 Chemical Transitions in Diffuse Clouds ...... 3
1.3 CO Dark Gas ...... 5
1.3.1 Hydroxyl (OH) ...... 7
1.3.2 CH emission ...... 8
1.3.3 H I Narrow Self Absorption ...... 9
1.4 Observations and Models ...... 9
1.5 Objectives ...... 13
2 Observations and Data Reduction 15
2.1 Target Stars ...... 15
2.2 Observations ...... 16
2.2.1 McDonald Observations ...... 20
ix 2.2.2 ESO Observations ...... 21
2.2.3 FUSE Observations ...... 22
2.2.4 HST Observations ...... 24
2.2.5 GOT C+ Observations ...... 27
2.3 Data Reduction ...... 30
2.3.1 McDonald Observatory Data Reduction ...... 30
2.3.2 ESO Data Reduction ...... 33
2.3.3 Known FUSE Detector Effects ...... 33
2.3.4 FUSE Data Reduction ...... 35
2.3.5 HST Data Reduction ...... 36
2.3.6 GOT C+ Reduction ...... 36
3 Component Fitting 37
3.1 ISMOD ...... 38
3.2 Fitting Ground-based Spectra ...... 39
3.2.1 Comparison with Literature Values ...... 76
3.2.2 HD 197770 ...... 86
3.3 Fitting of the H2 Bands ...... 89 3.3.1 b-Values ...... 90
3.3.2 Velocity Component Structure ...... 90
3.3.3 Additional Constraints ...... 91
3.3.4 H2 Results ...... 92 3.4 Fitting of the CO Bands ...... 98
3.4.1 Velocity Structure ...... 98
3.4.2 12CO Bands from HST ...... 100
3.4.3 12CO Bands from FUSE ...... 101
3.4.4 13CO Detections with HST ...... 106
x 3.4.5 12CO and 13CO Results ...... 106
3.5 Fitting of Additional Species ...... 113
3.5.1 C I ...... 113
3.5.2 C2 ...... 115 3.5.3 Species for the GOT C+ Project ...... 119
4 Analysis 131
4.1 Column Density Trends ...... 131
4.2 Excitation and Density ...... 133
4.2.1 H2 Excitation Temperatures ...... 134 4.2.2 CO Excitation Temperatures ...... 138
4.2.3 C I Excitation Analysis ...... 149
4.3 Chemistry and Density ...... 154
4.3.1 CN Chemistry ...... 154
4.3.2 CH+ like CH Chemistry ...... 172
4.4 h and χ Persei ...... 174
4.4.1 CO versus H2 ...... 174 4.4.2 The CN/CH+ and CH/CH+ Ratios ...... 178
4.4.3 Discussion ...... 183
4.5 The 12CO/13CO Isotopic Ratio ...... 184
4.5.1 Regimes ...... 185
4.5.2 Density Histograms for 12CO/13CO Sample ...... 187
4.6 Kinematic Associations of GOT C+ Species ...... 193
4.6.1 A Component by Component Discussion
for G014.8-1.0 ...... 193
4.6.2 Summaries for Other Pointings ...... 195
4.6.3 Updated Parallax Distances ...... 201
xi 5 Discussion 202
5.1 Combined Chemistry Picture ...... 202
5.2 GOT C+ Species Associations ...... 213
5.2.1 H I and Ca II Detections ...... 213
5.2.2 C+, Ca I,K I, CH+, CH, and CN Detections ...... 213
5.2.3 CO Detections ...... 214
5.2.4 Combined Absorption and Emission Picture ...... 216
5.3 CO emission and CO Dark Gas ...... 217
5.3.1 Chamaeleon Region ...... 217
5.3.2 Cepheus OB2 Region ...... 222
5.3.3 Perseus B5 Region ...... 224
5.3.4 Comparisons with Other Tracers ...... 227
5.4 Small Scale Structure toward h and χ Persei ...... 231
6 Conclusions 237
A Spectra from Ground-based Observations 261
BH2 Spectra from FUSE 289
C CO Spectra from FUSE 322
D CO Spectra from HST 349
E Spectra for GOT C+ Sight Lines 358
F Density and Temperature from C I 373
xii List of Tables
2.1 Stellar Data ...... 17
2.1 Continued ...... 18
2.1 Continued ...... 19
2.2 GOT C+ Pointings ...... 20
2.3 Observational Data for ESO UVES/VLT Targets ...... 23
2.4 Observational Data for FUSE Targets ...... 25
2.5 FUSE CO Bands ...... 26
2.6 FUSE Segments ...... 26
2.7 Observational Data for HST Targets ...... 27
2.8 Stellar Data for GOT C+ Targets ...... 29
3.1 Spectroscopic Data at Visible Wavelengths ...... 42
3.2 CH, CH+, and CN Results from McDonald Observatory Observations . 46
3.2 Continued ...... 47
3.2 Continued ...... 48
3.2 Continued ...... 49
3.2 Continued ...... 50
3.2 Continued ...... 51
3.2 Continued ...... 52
3.2 Continued ...... 53
3.2 Continued ...... 54
3.3 Ca II Results from McDonald Observatory Observations ...... 55
xiii 3.3 Continued ...... 56
3.3 Continued ...... 57
3.3 Continued ...... 58
3.3 Continued ...... 59
3.3 Continued ...... 60
3.3 Continued ...... 61
3.3 Continued ...... 62
3.4 Equivalent Width Measurements of CH, CH+, and CN from McDonald
Spectra ...... 63
3.4 Continued ...... 64
3.4 Continued ...... 65
3.4 Continued ...... 66
3.4 Continued ...... 67
3.4 Continued ...... 68
3.5 UVES Measurements for Southern Hemisphere Sight Lines ...... 69
3.5 Continued ...... 70
3.6 McDonald Results for the GOT C+ Project Sight Lines ...... 71
3.6 Continued ...... 72
3.6 Continued ...... 73
3.6 Continued ...... 74
3.6 Continued ...... 75
3.7 Total Column Density and Wλ Comparison ...... 78 3.7 Continued ...... 79
3.7 Continued ...... 80
3.7 Continued ...... 81
3.7 Continued ...... 82
3.7 Continued ...... 83
xiv 3.7 Continued ...... 84
3.7 Continued ...... 85
3.8 Total Column Density and Wλ Comparison: HD 197770 ...... 87 3.9 Component Column Density Comparison: HD 197770 ...... 88
3.10 H2 Column Densities ...... 96 3.10 Continued ...... 97
3.11 HST 12CO Bands ...... 100
3.12 HST 13CO Bands ...... 102
3.13 CO Results ...... 111
3.13 Continued ...... 112
3.14 C I Column Densities ...... 113
3.15 C I Component Structure ...... 114
3.16 Bands of C2 ...... 116
3.17 C2 Structure ...... 116
3.18 C2 Column Density ...... 116
3.19 Other C2 Results ...... 116 3.20 Velocities of Components with Corresponding Emission Features in the
Inner Galaxy ...... 124
3.20 Continued ...... 125
3.20 Continued ...... 126
3.20 Continued ...... 127
3.21 Velocities of Components with Corresponding Emission Features in the
Outer Galaxy ...... 128
3.21 Continued ...... 129
3.21 Continued ...... 130
4.1 H2 Excitation Temperatures ...... 136
xv 4.2 Excitation Temperatures and Density from T CO(1-0) in Sheffer et al. (2008)141 4.2 Continued ...... 142
4.3 CO Excitation Temperature Conversions ...... 143
4.4 Excitation Temperatures and Density from T CO(1-0) ...... 144
4.5 Excitation Temperatures and Density from T CO(2-1) & T CO(3-2) . . . . 145
4.6 C I Densities ...... 151
4.6 Continued ...... 152
4.7 τUV Determinations ...... 156 4.7 Continued ...... 157
4.8 CN Chemistry of Sight Lines in the Current Sample ...... 158
4.8 Continued ...... 159
4.8 Continued ...... 160
4.9 CN Chemistry of Sight Lines from the Sheffer et al. (2008) Sample . . . 161
4.9 Continued ...... 162
4.9 Continued ...... 163
4.9 Continued ...... 164
4.9 Continued ...... 165
4.9 Continued ...... 166
4.10 CN Chemistry of Additional Sight Lines from the Pan et al. (2004) Sample 167
4.10 Continued ...... 168
4.10 Continued ...... 169
4.11 CO and H2 Results for h and χ Percei ...... 182 4.12 13CO Results ...... 186
4.13 Comparison of Component Structure: G014.8-1.0 ...... 195
4.14 Comparison of Component Structure: G010.4+0.0 ...... 196
5.1 Comparison of Density Determinations (Current Data Set) ...... 206
xvi 5.1 Continued ...... 207
5.1 Continued ...... 208
5.2 Comparison of Different Density Determinations (previous efforts noted
in the text) ...... 209
5.2 Continued ...... 210
5.2 Continued ...... 211
5.2 Continued ...... 212
5.3 Summary of GOT C+ Associations ...... 215
5.4 Chamaeleon Data ...... 221
5.5 Data for Additional Tracers ...... 230
xvii List of Figures
3-1 HD 13841 McDonald spectra ...... 43
3-2 HD 197770 McDonald spectra ...... 44
3-3 HD 197770 CN spectra ...... 45
3-4 HD 108 FUSE H2 spectrum ...... 93
3-5 HD 170740 FUSE H2 spectrum ...... 94
3-6 HD 192641 FUSE H2 spectrum ...... 95 3-7 HD 108 HST CO spectra ...... 99
3-8 HD 108 HST CO spectra ...... 103
3-9 HD 108 FUSE CO spectra ...... 104
3-10 HD 25443 FUSE CO spectra ...... 105
3-11 HD 170740 HST 13CO spectra ...... 107
3-12 HD 108 HST d-X (5-0) CO spectrum ...... 108
3-13 HD 108 HST 13CO spectra ...... 109
3-14 All HD 46223 HST 13CO spectra ...... 110
3-15 HD 197770 C2 A-X (3-0) spectra ...... 117
3-16 HD 108927 C2 A-X (2-0) spectra ...... 118 3-17 HD 13841 McDonald spectra ...... 123
4-1 N (CO) versus N (H2) ...... 134
4-2 H2 Temperature Histogram ...... 137 4-3 Models of CO Excitation versus Density ...... 146
4-4 Density Histogram from CO Temperatures (J = 1-0) ...... 147
xviii 4-5 Density Histogram from CO Temperatures (J = 2-1) ...... 148
4-6 C I Density toward HD 108 ...... 153
4-7 CN Chemistry Pathways ...... 170
4-8 Density Histogram from CN chemistry ...... 173
4-9 N (CO) versus N (H2) ...... 175
4-10 N (CH) versus N (H2) ...... 176 4-11 N (CO) versus N (CH) ...... 177
4-12 CN/CH+ Ratio ...... 178
4-13 N (CH)/N (CH+) Histogram ...... 180
4-14 CN/CH+ Trend ...... 181
4-15 N (12CO)/N (13CO) versus N (12CO) ...... 188
12 12 13 4-16 Density Histogram from CO T J=1−0 for N ( CO)/N ( CO) Sample . . 190 4-17 N (12CO)/N (13CO) versus N (13CO) ...... 191
12 12 13 4-18 Density Histogram from CO T J=2−1 for N ( CO)/N ( CO) Sample . . 192
5-1 CO Emission Contours in Chamaeleon ...... 220
5-2 CO Emission Contours in Cep OB2 Region ...... 223
5-3 CO Emission Contours in Perseus B5 Region ...... 226
5-4 Ca II in h and χ Per ...... 234
5-5 CH in h and χ Per ...... 235
5-6 CH+ in h and χ Per ...... 236
6-1 The Structure of Interstellar Gas ...... 242
A-1 HD 108 McDonald spectra ...... 262
A-2 HD 5689 McDonald spectra ...... 263
A-3 HD 12882 McDonald spectra ...... 264
A-4 HD 13969 McDonald spectra ...... 265
A-5 HD 14053 McDonald spectra ...... 266
xix A-6 BD+56◦ 0501 McDonald spectra ...... 267
A-7 BD+56◦ 0563 McDonald spectra ...... 268
A-8 HD 14443 McDonald spectra ...... 269
A-9 HD 14476 McDonald spectra ...... 270
A-10 BD+56◦ 0578 McDonald spectra ...... 271
A-11 HD 14947 McDonald spectra ...... 272
A-12 HD 15629 McDonald spectra ...... 273
A-13 HD 16691 McDonald spectra ...... 274
A-14 HD 17505 McDonald spectra ...... 275
A-15 HD 17520 McDonald spectra ...... 276
A-16 BD+60◦ 0586 McDonald spectra ...... 277
A-17 HD 19243 McDonald spectra ...... 278
A-18 HD 19820 McDonald spectra ...... 279
A-19 HD 25443 McDonald spectra ...... 280
A-20 HD 25638 McDonald spectra ...... 281
A-21 HD 45314 McDonald spectra ...... 282
A-22 HD 46223 McDonald spectra ...... 283
A-23 HD 192303 McDonald spectra ...... 284
A-24 HD 192641 McDonald spectra ...... 285
A-25 HD 193793 McDonald spectra ...... 286
A-26 HD 214419 McDonald spectra ...... 287
A-27 HD 217086 McDonald spectra ...... 288
B-1 HD 5689 FUSE H2 spectrum ...... 290
B-2 HD 12882 FUSE H2 spectrum ...... 291
B-3 HD 13841 FUSE H2 spectrum ...... 292
B-4 HD 13969 FUSE H2 spectrum ...... 293
xx ◦ B-5 BD+56 0501 FUSE H2 spectrum ...... 294
◦ B-6 BD+56 0563 FUSE H2 spectrum ...... 295
B-7 HD 14053 FUSE H2 spectrum ...... 296
B-8 HD 14443 FUSE H2 spectrum ...... 297
B-9 HD 14476 FUSE H2 spectrum ...... 298
◦ B-10 BD+56 0578 FUSE H2 spectrum ...... 299
B-11 HD 14947 FUSE H2 spectrum ...... 300
B-12 HD 15629 FUSE H2 spectrum ...... 301
B-13 HD 16691 FUSE H2 spectrum ...... 302
B-14 HD 17505 FUSE H2 spectrum ...... 303
B-15 HD 17520 FUSE H2 spectrum ...... 304
◦ B-16 BD+60 0586 FUSE H2 spectrum ...... 305
B-17 HD 19243 FUSE H2 spectrum ...... 306
B-18 HD 19820 FUSE H2 spectrum ...... 307
B-19 HD 25443 FUSE H2 spectrum ...... 308
B-20 HD 25638 FUSE H2 spectrum ...... 309
B-21 HD 45314 FUSE H2 spectrum ...... 310
B-22 HD 46223 FUSE H2 spectrum ...... 311
B-23 HD 192303 FUSE H2 spectrum ...... 312
B-24 HD 193793 FUSE H2 spectrum ...... 313
B-25 HD 214419 FUSE H2 spectrum ...... 314
B-26 HD 217086 FUSE H2 spectrum ...... 315
B-27 HD 92964 FUSE H2 spectrum ...... 316
B-28 HD 97253 FUSE H2 spectrum ...... 317
B-29 HD 108927 FUSE H2 spectrum ...... 318
B-30 HD 149404 FUSE H2 spectrum ...... 319
B-31 HD 151932 FUSE H2 spectrum ...... 320
xxi B-32 HD 168076 FUSE H2 spectrum ...... 321
C-1 HD 5689 FUSE CO spectra ...... 323
C-2 HD 12882 FUSE CO spectra ...... 324
C-3 HD 13841 FUSE CO spectra ...... 325
C-4 HD 13969 FUSE CO spectra ...... 326
C-5 BD+56◦ 0501 FUSE CO spectra ...... 327
C-6 HD 14476 FUSE CO spectra ...... 328
C-7 HD 15629 FUSE CO spectra ...... 329
C-8 HD 16691 FUSE CO spectra ...... 330
C-9 HD 17505 FUSE CO spectra ...... 331
C-10 HD 17520 FUSE CO spectra ...... 332
C-11 BD+60◦ 0586 FUSE CO spectra ...... 333
C-12 HD 19243 FUSE CO spectra ...... 334
C-13 HD 19820 FUSE CO spectra ...... 335
C-14 HD 25638 FUSE CO spectra ...... 336
C-15 HD 45314 FUSE CO spectra ...... 337
C-16 HD 46223 FUSE CO spectra ...... 338
C-17 HD 97253 FUSE CO spectra ...... 339
C-18 HD 149404 FUSE CO spectra ...... 340
C-19 HD 151932 FUSE CO spectra ...... 341
C-20 HD 168076 FUSE CO spectra ...... 342
C-21 HD 170740 FUSE CO spectra ...... 343
C-22 HD 14053 and HD 14947 FUSE CO spectra...... 344
C-23 HD 193793 and HD 197770 FUSE CO spectra...... 345
C-24 HD 214419 and HD 217086 FUSE CO spectra...... 346
C-25 BD+56◦ 0563 and BD+56◦ 0578 FUSE CO spectra...... 347
xxii C-26 HD 14443, HD 92964, HD 192303, HD 192641 FUSE CO spectra. . . . 348
D-1 HD 13841 HST CO spectra ...... 350
D-2 HD 25443 HST CO spectra ...... 351
D-3 HD 108927 HST 13CO spectra ...... 352
D-4 HD 46223 HST CO spectra ...... 353
D-5 HD 46223 HST d-X (5-0) CO spectra ...... 354
D-6 HD 46223 HST 13CO spectra ...... 355
D-7 HD 108927 HST CO spectra ...... 356
D-8 HD 170740 HST CO spectra ...... 357
E-1 HD 168607 spectra ...... 359
E-2 HD 165918 spectra ...... 360
E-3 HD 167498 spectra ...... 361
E-4 HD 167812 spectra ...... 362
E-5 HD 169754 spectra ...... 363
E-6 HD 174509 spectra ...... 364
E-7 BD+49◦ 3482 spectra ...... 365
E-8 BD+49◦ 3484 spectra ...... 366
E-9 HD 240179 spectra ...... 367
E-10 HD 240183 spectra ...... 368
E-11 HD 47073 spectra ...... 369
E-12 HD 260737 spectra ...... 370
E-13 HD 55469 spectra ...... 371
E-14 HD 55981 spectra ...... 372
F-1 C I Density toward HD 13841 ...... 374
F-2 C I Density toward HD 25443 ...... 375
F-3 C I Density toward HD 46223 ...... 376
xxiii F-4 C I Density toward HD 108927 ...... 377
F-5 C I Density toward HD 170740 ...... 378
xxiv Chapter 1
Introduction
The focus of this Dissertation is the transition between diffuse atomic clouds and dense molecular clouds. The first section of this chapter provides a brief review of the relevant aspects of the interstellar medium (ISM) and different types of clouds. This is followed by a discussion of the different chemical transitions seen in interstellar clouds. The next section introduces CO dark gas and some of the probes that have been used to analyze it. Section 1.4 presents the different pictures of diffuse clouds and their tracers that come from absorption and emission observations. The final section (Section 1.5) gives the overall objectives of this Dissertation.
1.1 ISM Overview
The interstellar medium is made of gas and dust. Most of the mass of the ISM is found in the gas phase (≈99%). Hydrogen makes up about ≈90% of the nucleons in the gas phase of the ISM, while helium makes up about ≈10% and heavier elements like carbon, nitrogen, and oxygen constitute less than ≈1% of the nucleons (Snow &
McCall 2006).
Although dust only makes up about 1% of the total mass of the ISM, the dust grains are a catalyst for interactions between particles and has a large impact on the chemistry of gas in the ISM. H2 is produced on grains by surface processes; H atoms 1 stick to the grain’s surface, and then tunnel or migrate between binding sites to react
and form H2, which is eventually ejected back into the gas (De Becker 2013). Gas
phase pathways for H2 formation are much slower (Snow & McCall 2006). To be observable in the ISM, molecules must form fast enough to balance their destruction.
In the diffuse ISM, most molecules are in their ground electronic and vibrational
states. The excitation of the different rotational states is controlled by radiative and
collisional processes (Snow & McCall 2006). Molecules with dipole moments have
allowed rotation transitions and excitation that can occur through stimulated ab-
sorption, while de-excitation can occur through spontaneous and stimulated emission
(Snow & McCall 2006). Molecules without electric dipole moments can still have
magnetic or higher order electric moments and can be excited by photon absorption
in electronic transitions (Snow & McCall 2006). Once the relevant processes are un-
derstood, the populations of the excited levels provide information about the physical
conditions, such as the density and kinetic temperature of the gas.
In addition to the formation of molecules on grains, the absorption and scatter-
ing of starlight by dust controls most of the attenuation of starlight passing through
interstellar clouds (Snow & McCall 2006). Of particular importance is the far ultra-
violet (UV) radiation from early type stars, which dominates the destruction of small
molecules in the ISM (De Becker 2013). The gas in diffuse atomic clouds is fully
exposed to starlight; as such, the molecular abundances are very low. Hydrogen is al-
most entirely in neutral atomic form, while carbon is almost fully ionized. Snow & Mc-
Call (2006) suggest using the molecular hydrogen fraction, 2N(H2)/[N(H I)+2N(H2)], to help distinguish different types of gas, where N (X) is the column density of the
species. Snow & McCall (2006) define gas with a low molecular hydrogen fraction
(≤ 5 to 10%) as diffuse atomic gas and a high molecular hydrogen fraction (≥ 90%) as dense molecular gas. Diffuse molecular gas and so called ‘translucent’ gas are in between the two extremes, with the fraction of elemental carbon in C+ differentiating
2 them. As the molecular hydrogen fraction changes, so do the species present and the dominant processes in the chemistry of the gas.
1.2 Chemical Transitions in Diffuse Clouds
The types of clouds can be described further using their different chemistries or chemical transitions. Diffuse atomic gas consists mainly of hydrogen atoms and can be traced by the H I hyperfine transition at 21 cm or through UV observations of
Lyman-alpha absorption. As the amount of molecular hydrogen increases, it begins to self shield. This process suppresses the wavelengths of the interstellar radiation field that would photodissociate the H2 deeper into the cloud, resulting in the transition from diffuse atomic gas to diffuse molecular gas (Hollenbach et al. 1971). However, the radiation is still strong enough to photodissociate most of the CO that forms at this stage (Wolfire et al. 2010).
Molecular hydrogen is a symmetric and homonuclear molecule without a perma- nent electric dipole moment. Molecular hydrogen was predicted to be abundant in models of diffuse clouds (Hollenbach et al. 1971; Glassgold & Langer 1974, 1976;
Black & Dalgarno 1976, 1977). In the diffuse ISM, the readily observable probes of
H2 are the UV electronic transitions in the Lyman and Werner bands below 1115 A˚
(Snow & McCall 2006; Spitzer & Zabriskie 1959). The first observations of strong H2 absorption were made by Carruthers (1970) with rocket based spectroscopy. Later, ro- tational excited absorption lines were observed with the Copernicus telescope (Spitzer
& Jenkins 1975), which launched in 1972.
The transition between diffuse molecular gas and translucent gas comes from the self shielding of CO, just as the transition between diffuse atomic gas and diffuse molecular gas involves H2 self shielding. Carbon monoxide is a stable molecule with a bonding energy of 11.09 eV. Since most photons with energies greater than 13.6
3 eV are consumed in the ionization of hydrogen atoms, the photodissociation of CO
occurs through photons with energies between 11.09 eV and 13.6 eV, corresponding
to wavelengths between 912 and 1118 A˚ (De Becker 2013). The electronic transitions
of CO can be observed in the UV. Unlike H2, CO has dipole allowed rotational transitions in the millimeter wave band and vibrational transitions in the infrared.
Radio observations by Wilson et al. (1970) and ultraviolet spectroscopy by Smith &
Stecher (1971) confirmed the existence of CO in the interstellar medium.
The first indication of a relationship between CO and H2 column densities was made by Federman et al. (1980), who found a slope of approximately two. Further
work (Federman et al. 1994; Liszt & Lucas 1998; Rachford et al. 2002; Pan et al.
2005; Burgh et al. 2007; Sonnentrucker et al. 2007; Sheffer et al. 2008) indicated
different slopes in the relationship between CO and H2 column densities for log[N (H2)] greater than or less than 20.4. This break in the slope comes from a change in the photochemistry. Models by Bally & Langer (1982) and van Dishoeck & Black (1988) showed the importance of line photodissociation and increased self shielding to CO photochemistry through molecular clouds. Their results mirror the observed change in the photochemistry by Sheffer et al. (2008). If the trend between CO and H2 is extended to higher column densities, Sheffer et al. (2008) showed it is consistent with the results from millimeter wave observations of dark clouds.
As extinction increases, carbon eventually becomes primarily molecular, moving into the regime of dense molecular gas where the molecular hydrogen fraction is ≥
90% and the electron abundance is very low, decreasing the impact of dissociative recombination and giving rise to a rich ion molecule chemistry (Snow & McCall 2006;
De Becker 2013). Although dense molecular gas is important with a rich chemistry, the focus here is on diffuse molecular and translucent gas. It is in these regions that models of interstellar clouds predict the majority of CO dark gas to reside (van
Dishoeck & Black 1988; Tielens 2005; Wolfire et al. 2010).
4 1.3 CO Dark Gas
In molecular clouds, H2 is primarily traced by emission from other more easily
observed species such as CO. The X CO (X CO = NH2 /WCO) factor is a linear conversion factor between the integrated CO(J = 1-0) line intensity and the H2 column density.
This XCO conversion factor is often approximated as constant but studies have shown cloud to cloud variations in diffuse molecular gas (Magnani & Onello 1995; Magnani et al. 1998) and variations at larger scales along metallicity and UV gradients (Abdo et al. 2010; Ackermann et al. 2011; Pineda et al. 2013; Bolatto et al. 2013). In addition, there is ample evidence that CO emission does not trace all the molecular gas in different environments equally (van Dishoeck & Black 1988; Magnani & Onello
1993; Wannier et al. 1993; Grenier et al. 2005; Douglas & Taylor 2007; Barriault et al. 2010; Liszt et al. 2010; Liszt 2011; Cotten & Magnani 2013). Magnani et al.
(2003) confirmed that CO is not a linear tracer of the molecular hydrogen column density throughout diffuse molecular clouds. A region of molecular gas, not observable through CO or H I emission, was first predicted to exist in diffuse molecular and translucent gas by van Dishoeck & Black (1988).
The far infrared observations of Paradis et al. (2012) and the γ ray observations of Grenier et al. (2005) showed the ubiquity of gas not traced by 21 cm H I or 2.6 mm
CO emission in the ISM. Planck Collaboration et al. (2015) used γ ray observations with the Fermi Large Area Telescope and thermal dust emission observed with Planck and the Infrared Astronomical Satellite (IRAS) to provide measures of the total gas.
The dust content can also be traced by extinction and reddening measurements. The
Dark Neutral Medium (DNM) shows up as γ ray and dust excesses over the N HI and
W CO expectations (Blitz et al. 1990; Reach et al. 1994, 1998; Magnani et al. 2003; Planck Collaboration et al. 2015). CO dark gas is the molecular part of the DNM.
Remy et al. (2018b) extends the results of Planck Collaboration et al. (2015) by
5 focusing on the decomposition of a small subset of 15 clouds. They find that the molecular mass in the DNM of diffuse clouds (with AV ≤ 1 mag) often exceeds the
H2 mass derived from CO observations. The work of Remy et al. (2017) indicates that changes in dust properties can lead to XCO from dust emission that are 30% to 130% larger than XCO from γ ray estimates. They find that the dust opacity increases with the density of the gas. This is due to changes in the dust emissivity (Remy et al.
2018a). Variations in dust properties have also been suggested by Reach et al. (2015) as a possible source of the variations seen in the gas to dust ratio. However, in later work, Reach et al. (2017a) determine that changes in dust properties cannot account for the range of the gas to dust ratio observed in diffuse clouds and instead attribute the variations to CO dark gas.
Typically, 21 cm H I emission is assumed to be optically thin. Fukui et al. (2014) and Fukui et al. (2015) argue that dark gas can be explained by incorrect deter- minations of H I opacities and find cold H I gas masses 2 to 2.5 times that of the optically thin calculations. However, Li et al. (2015) analyzed the Perseus Molecular
Cloud and found only a 10% increase in the H I gas mass from optically thick H I calculations, not enough to account for the majority of the CO dark gas. Li et al.
(2015) also find that the distributions of CO dark gas and optically thick H I do not tend to follow each other. The search for cold optically thick H I in interstellar clouds by Reach et al. (2017b) yielded only a small amount of the dark gas that can be accounted for by H I in cold atomic gas. They conclude that cold atomic gas is only a minor component of the total inferred dark gas. Recent work by Murray et al.
(2018) confirms that optically thick H I does not dominate the dark gas in the local
ISM, and they suggest the signatures of dark gas, seen as excess dust emission, are a combination of changes in dust emissivity and H2 in CO dark gas. Furthermore, Togi
& Smith (2016) find substantial CO dark gas at low metallicity using H2 emission, which should not suffer dust emission biases. As noted above, Reach et al. (2017a)
6 find that changes in dust properties cannot account for the range of the gas to dust
ratio seen in diffuse clouds and conclude that CO dark gas is likely composed of H2. CO dark gas is ‘dark’ in CO emission because of its low abundance and the low kinetic temperatures in diffuse gas (T ≤ 100 K) giving rise to very low fluxes from
rotational H2 line emission (Wolfire et al. 2010). Although the gas is seen as ‘dark’ in CO emission, it does emit in the C+ fine-structure line (Langer et al. 2010; Pineda
et al. 2010; Liszt 2011; Pineda et al. 2013). There are several other molecular and
atomic species, other than CO, that can be used to trace or infer properties of the
diffuse gas found in the envelope around molecular clouds, where we expect to find
CO dark gas. In particular, there are the OH ground state lines (1665/1667 MHz)
(Wouterloot 1981; Wannier et al. 1993; Liszt & Lucas 1996), CH ground state line
emission (3264, 3335, and 3349 MHz) (Magnani et al. 1989; Xu & Li 2016), the fine-
structure lines of atomic carbon (especially 492 GHz) (Ingalls et al. 1997), and C+ line emission (158 µm) (Langer et al. 2010; Pineda et al. 2010; Liszt 2011; Pineda et al. 2013; Langer et al. 2014). The following subsections briefly describe some of these tracers that are not discussed in depth elsewhere.
1.3.1 Hydroxyl (OH)
OH is expected to be abundant in diffuse interstellar molecular gas because it forms along with H2 under similar conditions (Li et al. 2018). The two main 18 cm lines of OH were first detected in absorption in the ISM by Weinreb et al. (1963) and in emission by Weaver et al. (1965). Diffuse OH emission is typically about 100 times fainter than the H I 21 cm line (Dickey et al. 2013). Lewis et al. (1988) found that clouds with CO emission contain measurable amounts of OH that falls off slower than CO, and extends beyond the range of CO emission (Wannier et al. 1993; Liszt
& Lucas 1996). Observations by Barriault et al. (2010) and Allen et al. (2012) show there is little correlation between OH column density and WCO. Since OH forms
7 along with H2, this implies that large amounts of H2 gas are not being detected in surveys of CO emission.
Li et al. (2015) conducted a followup CO survey of 79 Millennium sight lines with
OH absorption and found OH to be a good tracer of dark molecular gas, co-existing
with atomic and molecular hydrogen in regions without CO emission. Sancisi et al.
(1974) found that H I, dust, OH, and CH are physically associated in the Perseus
OB2 cloud. Li et al. (2015, 2018) claim OH roughly resides between the self shielding
13 thresholds for H2 and CO (0.05 ≥ AV ≥ 2). Porras et al. (2014) found that the association of near UV OH absorption with gas rich in CN is attributed to the need for high enough density and molecular fraction before detectable amounts are seen.
While OH+ leads to OH production, chemical arguments suggest that their abun- dances are controlled by different sets of conditions and that they coexist with different sets of observed species (Porras et al. 2014). Absorption from OH+ was first detected at near ultraviolet wavelengths through measurements acquired with the VLT/UVES instrument (Kre lowski et al. 2010). Kre lowski et al. (2010) suggested that OH+ ab- sorption traces CH+ the best among the molecular species seen at visible wavelengths.
Porras et al. (2014) confirmed this connection for individual velocity components or clouds.
1.3.2 CH emission
The 3335 MHz hyperfine ground state transition of CH was first observed in diffuse clouds by Lang & Willson (1978) and Willson (1981). The CH emission line has been found to be well correlated with E(B − V ) in diffuse clouds but not with CO(J =1-0) emission (Magnani et al. 2003), indicating that CH emission continues to trace the total gas beyond the range accessible with CO emission.
Federman (1982) first showed N (CH) to be linearly related to N (H2). Additional observation confirmed a tight relation between N (CH) and N (H2) (with a slope of
8 ≈0.97) in diffuse and dense molecular clouds (Danks et al. 1984; Rachford et al. 2002;
Pan et al. 2005; Sheffer et al. 2008). The broad wings of the CH 3335 MHz line
are not seen in the CO(J = 1-0) line (Magnani & Onello 1993). Magnani & Onello
found that the broad wings of CH provide a way to trace the low density molecular
hydrogen gas between the dense molecular gas and the diffuse atomic gas envelope.
Furthermore, Xu & Li (2016) found CH emission to be a better tracer of the total
column density than OH at AV between 0.8 and 2.1 mag. This magnitude range roughly agrees with the range expected for diffuse molecular clouds.
1.3.3 H I Narrow Self Absorption
The absorption of radiation from a distant source by cool H I gas is an important tool for studying atomic hydrogen in interstellar clouds. If the radiation from the source is H I emission, the absorption from the foreground cloud produces a H I narrow self absorption (HINSA) feature (Li & Goldsmith 2003; Goldsmith & Li 2005;
Goldsmith et al. 2007). A significant amount of H I continues to exist inside molecular clouds due to the destruction of H2 by penetrating cosmic rays (Li & Goldsmith 2003). Li & Goldsmith (2003) found that the atomic hydrogen in HINSA features was associated with the molecular material traced by emission from OH and CO isotopologues in dark clouds. Gas with HINSA features are also observed to follow the distribution seen in emission from CO isotopologues (Goldsmith & Li 2005).
1.4 Observations and Models
The physical conditions in interstellar clouds range from very diffuse conditions in atomic gas to those in dense molecular clouds in which star formation takes place.
The atomic hydrogen in diffuse atomic gas transitions to molecular hydrogen seen in diffuse molecular gas, which contains a mix of atomic and molecular gas. While
9 atomic hydrogen can be observed at 21 cm in emission in this material or with Lyman-
alpha absorption, molecular hydrogen is usually observed in absorption against a
background source. Many times, species observed in absorption are only combined
with other species in absorption, while species observed in emission are combined
with other species in emission, creating two separate views of the same environment.
By combining the information from different species, one can trace a larger range
of physical conditions, and when considered as an ensemble, they provide a more
complete picture of an entire interstellar cloud.
Studies connecting UV/visible absorption and long-wavelength emission began
more 40 years ago. Knapp & Jura (1976) sought CO emission from gas observed
in absorption with the Copernicus satellite. Liszt (1979) acquired data on CH and
CO(J = 1 → 0) emission along the line of sight to ζ Oph, finding material with a very small thermal line width. Subsequent observations of CO(J = 2 → 1) emission
(Crutcher & Federman 1987) and absorption from CH and CN toward the star (Lam- bert et al. 1990; Crawford et al. 1994) confirmed the component structure seen by
Liszt (1979). Willson (1981) examined the correspondence between CH emission and absorption toward background stars. Building on this correspondence, Federman &
Willson (1982) used CH measurements to suggest an association between dark clouds and their envelopes of diffuse molecular gas, while Mattila (1986) showed more clearly the connection between diffuse molecular gas and dark clouds. For the material in the direction toward HD 29647, Crutcher (1985) combined data at radio and visible wavelengths for a comprehensive analysis of the material. Maps of molecular emission in the vicinity of background stars were obtained to learn more about the association of diffuse and dark molecular gas (Federman et al. 1984; Gredel et al. 1992, 1994).
Diffuse atomic and molecular gas have been traditionally studied using absorption in the UV and visible portions of the spectrum. For instance, CN absorption traces
−3 relatively dense gas (nH ∼ 300 cm ) (Pan et al. 2005; Sheffer et al. 2008; Rice et al.
10 2018). The density (nH) refers to total proton density, n(H I)+2n(H2). In the portions of the cloud surrounding this gas without detectable amounts of CN absorption, we
find absorption from CH, CH+, K I, Ca I, and Ca II as the density decreases. Among
these species, the CN profile is usually much less complicated than profiles of other
species for a given sight line because it is only probing the denser material (Pan
et al. 2005). The column densities of CN components are correlated with column
densities of corresponding CH components surrounding the denser CN regions (Pan
et al. 2005). Welty & Hobbs (2001) found an essentially linear relationship between
N (K I) and N (CH) by using total column densities along lines of sight. The column
densities N (Ca I) and N (K I) are also well correlated (Welty et al. 2003). Figure
6 from Pan et al. (2005) shows a schematic of this idea, indicating the overlapping
layered locations of tracers at visible wavelengths within a cloud.
−3 • Ca II samples low density atomic gas (nH ≤ 10 cm )
+ • K I, Ca II, CH , and CH sample somewhat denser molecular gas (nH ≈ 100 cm−3)
• CN and CO sample the densest regions of diffuse molecular clouds (nH ≈ 300 cm−3)
While absorption spectra are sensitive to low densities and column densities, emis- sion spectra are sensitive to density via the collisional rate and are more difficult to detect in low density environments. Recently, it has become possible to study these environments in emission from the 158 µm fine-structure line of C+ (Langer et al.
2010; Pineda et al. 2010; Liszt 2011; Pineda et al. 2013). However, because C+ has only one fine-structure line, it cannot completely characterize the emitting gas with- out making assumptions about the physical conditions (Velusamy et al. 2010; Langer et al. 2014). This Dissertation is meant to complement the Herschel Space Obser- vatory key program GOT C+ (Galactic Observations of Terahertz C+), which was 11 2 2 designed to study the ISM by connecting the emission from the [C II] P3/2 → P1/2
fine-structure line at 158 µm with emission from H I at 21 cm and CO at 2.6 mm.
The GOT C+ project is described in detail by Langer et al. (2010). The analyses
of Langer et al. (2014) were based on tracers (H I, 12CO, 13CO) observed at radio
and millimeter wavelengths. In order to discern the types of interstellar environment
producing C+ emission, they adopted the following picture.
• Diffuse Atomic Clouds (Warm Neutral Medium): Components seen in H I only;
the densities and column densities are too low to produce C+ emission above
−3 the GOT C+ sensitivity limit. (nH ≤ 1 cm )
• Diffuse Atomic Clouds (Cold Neutral Medium): Components seen in H I, and/or
C+, with 12CO and 13CO emission well below the detection limit, with the
assumption that this H I is associated with the envelope around denser colder
gas. Here, some of the C+ emission comes from H I and some from CO dark
−3 −3 gas. (1 cm ≤ nH ≤ 300 cm )
• CO Dark Gas: Components seen in H I, and/or C+, but with no 12CO and 13CO
−3 emission above the detection limit. (nH ≈ 300 cm )
• Molecular Clouds: Components seen in H I, and/or C+, 12CO, and 13CO emis-
13 −3 sion, with the distinction that clouds with CO are denser. (300 cm ≤ nH ≤ 1000 cm−3)
Studies of atomic and molecular absorption and emission in the diffuse ISM have their own seemingly separate paradigms. However, measurements of atomic and molecular absorption at visible and UV wavelengths provide complementary data on some of the environments seen in the GOT C+ survey. By associating probes seen at visible and UV wavelengths with the results from the GOT C+ survey, I can
12 pursue an integrated study of the diffuse ISM, incorporating data from radio to UV
wavelengths.
1.5 Objectives
The first objective of the Dissertation is to understand the physical conditions
of the changing environment between diffuse and dense gas. To this end, we utilize
the column densities and excitation temperatures from CO and H2 to determine the gas density. The resulting gas densities from this method are compared to densities
inferred from other methods such as CN chemistry or C2 and C I excitation. The densities allow us to interpret the trends from the combined set of tracers.
A second objective of the Dissertation is to combine the emission and absorption
data to begin bridging the gap between two different views of the diffuse interstel-
lar medium: the paradigm created by tracers seen in absorption at visible and UV
wavelengths and the one created by tracers seen in emission at longer wavelengths.
The third objective is to investigate how the different methods and tracers relate
to the location and physical conditions in CO dark gas. Archival maps of molecular
emission in the vicinity of background stars were augmented with additional data to
learn more about the association of diffuse and dark molecular gas. Groupings of
sight lines, such as those toward h and χ Persei or Chamaeleon, provide a chance for further characterization of the environment. The Chamaeleon Region in particular helps illuminate CO dark gas.
The rest of the Dissertation is structured as follows. Chapter 2 presents the observations and data reduction. The fitting of lines from the observed species in the different spectral regions is described in Chapter 3. Chapter 4 illustrates how the data are used to determine the physical conditions in the gas (density, temperature, structure). Chapter 5 presents the combined picture from the different tracers and
13 their application to CO dark gas. Chapter 6 summarizes the results and discusses future extensions of this work.
14 Chapter 2
Observations and Data Reduction
2.1 Target Stars
High resolution observations of absorption from interstellar gas require the use of a bright, ideally, rapidly rotating star, as the light source; the sharp interstellar absorption lines are superposed on broader stellar features. Typically, there are mul- tiple interstellar clouds (or cloud clumps) along a given line of sight whose existence is revealed by absorption at distinct Doppler velocities. The background targets in our survey had to be relatively bright (B, V ≤ 10) O, B and A stars, see Table 2.1 for individual details. Additionally, only targets with available UV data (for the analysis of CO and H2) were considered for our primary sample. The stellar data for the 36 stars in our primary sample are presented in Table
2.1. The stars are listed by their entry in the Henry Draper (HD) catalogue, in order of increasing right ascension (RA), followed by Declination (Dec), B, and V magni- tudes, and spectral type which were all taken from the SIMBAD Database, operated at Centre de Donn´eesAstronomiques de Strasbourg (CDS), France (Wenger et al.
2000). References for spectral types, B, and V are provided in the footnotes. The total exposure time and distances are also provided. The distances used through- out this dissertation were determined by spectroscopic parallax and are known to
15 approximately 20% based on average uncertainties in magnitudes (∼0.02). Recently,
Bailer-Jones et al. (2018) performed a reanalysis of the GAIA DR2 parallax distances
(Gaia Collaboration et al. 2017). These parallax distances are listed in Table 2.1.
Section 4.6.3 discusses the impacts of the revised distances.
A complimentary sample of seventeen sight lines were chosen by examining the
SIMBAD database (Wenger et al. 2000); only GOT C+ pointings accessible to the
Northern Hemisphere were considered. The background targets in this sample also
had to be relatively bright B and A stars with amounts of extinction per kpc typical
for directions with diffuse molecular gas. Further details are provided in the Section
2.2.5. In addition, the lines of sight had to be within 30 arcminutes of a Herschel
pointing to have relevant emission data. These are listed in Table 2.2.
2.2 Observations
In order to study cloud structure and physical conditions of the clouds, 36 stars
in the primary sample were observed for interstellar CN, CH, CH+, Ca I, Ca II, and
K I absorption in the visible. K I was only observed for a subset of the sight lines from December of 2012. The high resolution visible spectra were acquired at two sites. Northern stars were observed with the 2.7 m Harlan J. Smith Telescope at the McDonald Observatory of the University of Texas in 2010 and 2012 with the echelle spectrograph (Tull et al. 1995). Archival data for southern stars observed at the European Southern Observatory (ESO) VLT/UVES were also used. These observations were combined with archival ultraviolet spectra from the Far Ultraviolet
Spectroscopic Explorer (FUSE) observations, and a few additional supplementary observations from Hubble Space Telescope (HST ).
16 Table 2.1: Stellar Data
Star RA (J2000) DEC (J2000) BV Spectral Distancea Distanceb Exp Timec (Mag) (Mag) Type (pc) (pc) (sec) Northern Targets HD 108 00 06 03.39 +63 40 46.76 7.491 7.381 O4-8f2 1800 2400 7200 HD 5689 00 59 47.59 +63 36 28.24 9.413 9.193 O63 3300 2950 14400 HD 12882 02 08 45.44 +65 02 14.72 7.924 7.624 B6Ia5 3500 900 14400 HD 13841 02 16 46.39 +57 01 45.66 7.541 7.391 B2Ib6 3300 3000 7200 HD 13969 02 17 49.85 +57 05 25.51 9.131 8.957 B1IV8 2500 2450 7200 1 7 8 17 HD 14053 02 18 23.05 +57 00 36.68 8.70 8.51 B1II 3200 2600 9000 BD+ 56◦ 0501 02 18 29.83 +57 09 03.14 9.631 9.357 B0.5I8 7900 2100 10800 BD+ 56◦ 0563 02 21 43.39 +57 07 32.84 9.771 9.457 B1IIIe9 3300 2750 10800 HD 14443 02 22 00.59 +57 08 41.87 8.391 8.057 BC2Ib5 3700 2650 5400 HD 14476 02 22 16.96 +57 16 18.94 9.131 8.857 B0.5III8 3300 2350 5400 BD+ 56◦ 0578 02 22 17.69 +57 07 24.67 9.571 9.257 B2III10 2400 3450 9000 HD 14947 02 26 46.99 +58 52 33.12 8.401 8.051 O5If11 1700 3300 5400 HD 15629 02 33 20.59 +61 31 18.18 8.831 8.401 O5e11 2500 2050 5400 HD 16691 02 42 52.03 +56 54 16.47 9.0712 8.7012 O4If11 2400 2150 3600 HD 17505 02 51 07.98 +60 25 03.87 7.491 7.101 O6Ve11 1200 2550 3600 HD 17520 02 51 14.46 +60 23 09.82 8.531 8.291 O9V11 1800 1900 5400 BD+ 60◦ 0586 02 54 10.67 +60 39 03.50 8.7012 8.5012 O8III13 2900 2750 3600 Table 2.1: —Continued
Star RA (J2000) DEC (J2000) BV Spectral Distancea Distanceb Exp Timec (Mag) (Mag) Type (pc) (pc) (sec) HD 19243 03 08 54.18 +62 23 04.55 6.714 6.504 B1V12 790 950 3360 HD 19820 03 14 05.33 +59 33 48.48 7.541 7.151 O8.5III11 1600 1300 12600 HD 25443 04 06 08.06 +62 06 06.60 6.994 6.784 B0.5III3 1000 1600 5400 HD 25638 04 07 49.29 +62 19 58.58 7.263 6.933 B0III14 1000 1100 3600 HD 45314 06 27 15.78 +14 53 21.22 6.691 6.601 O9:pe11 670 800 2700 HD 46223 06 32 09.31 +04 49 24.70 7.451 7.321 O5e11 1900 1650 3600 HD 192303 20 12 47.80 +38 13 39.60 9.201 8.941 B1III15 2500 2000 9000 HD 192641 20 14 31.77 +36 39 39.60 8.221 7.951 WCp16 1300 2100 6780 HD 193793 20 20 27.98 +43 51 16.28 7.281 6.891 WC7p+O517 480 16500 7200 d 4 4 3 18 HD 197770 20 43 13.68 +57 06 50.39 6.59 6.34 B2III 640 900 3600 HD 214419 22 36 53.95 +56 54 20.99 9.201 8.871 WN18 2800 3000 9000 HD 217086 22 56 47.19 +62 43 37.65 8.191 7.711 O7Vn19 960 850 10800 Table 2.1: —Continued
Star RA (J2000) DEC (J2000) BV Spectral Distancea Distanceb Exp Timec (Mag) (Mag) Type (pc) (pc) (sec) Southern Targets HD 92964 10 42 40.57 -59 12 56.72 5.591 5.401 B2.5Ia20 1200 3900 10170 HD 97253 11 10 42.04 -60 23 04.27 7.261 7.111 O5IIIe19 1900 3550 1577 HD 108927 12 32 20.40 -78 11 37.00 7.834 7.764 B5v21 330 340 3923 HD 149404 16 36 22.56 -42 51 31.90 5.8722 5.4722 O9Ia19 880 1350 5760 HD 151932 16 52 19.25 -41 51 16.26 6.7222 6.4922 WN7h23 2600 1250 8950 1 1 11 19 HD 168076 18 18 36.43 -13 48 02.03 8.64 8.23 O3.5Vf+O7.5V 2400 1550 7472 HD 170740 18 31 25.69 -10 47 45.01 5.931 5.761 B2V24 330 230 30220 aSpectroscopic distance. bParallax distance from GAIA DR2 using Bailer-Jones et al. (2018). cExp time refers to the total exposure time of observations. dHD 197770 has additional observations with the Hobby-Eberly Telescope. 1-Ducati (2002), 2-Martins et al. (2012), 3-Reed (2003), 4-Høg et al. (2000a), 5-Walborn (1971), 6-Lennon et al. (1992), 7-Cloutier et al. (2014), 8-Currie et al. (2010), 9-Currie et al. (2008), 10-Saesen et al. (2010), 11-Sota et al. (2011), 12-Zdanaviˇcius& Zdanaviˇcius(2002), 13-Mathys (1989), 14-Lutz & Lutz (1977), 15-Barbier (1962), 16-Williams et al. (2001), 17-Shenavrin et al. (2011), 18-Hamann et al. (2006), 19-Sota et al. (2014), 20-Hiltner et al. (1969), 21-Houk & Cowley (1975), 22-Zacharias et al. (2012), 23-Houk (1978), 24-Houk & Swift (1999) Table 2.2: GOT C+ Pointings
ID l b G014.8-1.0 14.7826 -1.0 G010.4+0.0 10.4348 0.0 G010.4+0.0 10.4348 0.0 G015.7+1.0 15.6522 1.0 G020.0+0.0 20.0000 0.0 G032.6+0.0 32.5532 0.0 G091.7+1.0 91.6981 1.0 G109.8+0.0 109.8113 0.0 G207.2-1.0 207.1700 -1.0 G225.3+0.0 225.2830 0.0
Additional observations at visible wavelengths for the GOT C+ sight lines were
also taken at the McDonald Observatory in December of 2012, July of 2014, and
October of 2017. Section 2.2.5 discusses the stellar data and observational details for
these lines of sight.
2.2.1 McDonald Observations
Twenty nine of the sight lines for the primary sample were observed with the 2d
coude spectrograph (cs21) (Tull et al. 1995) at the McDonald Observatory. The 79
gr mm−1 grating E1 and the TK3 charge-coupled device (CCD) were used for the
observations. A cross-dispersed echelle spectrometer and the availability of a 2048 x
2048 CCD made it possible to get CN λ3874, CH λ4300, CH+ λ4232, Ca I λ4226, and
Ca II λ3933 absorption features in a single exposure by setting the cross disperser to its 2nd order, and centering the 56th order of the echelle grating spectrum to 4065
A.˚ K I λ7699 spectra were obtained by setting the cross disperser to its 1st order and centering the 31st order spectrum to 7165 A.˚ Slit 2 with a width of 145 µm was used
for all observations.
An individual exposure time of 30 minutes was chosen to limit the effect of cosmic
20 rays during the data collection process. Cosmic rays create blemishes in the data when they strike the detector. These blemishes can be removed if there are only a few present in an exposure, but the difficulty involved in removing them increases as their population increases.
A variety of auxiliary exposures were taken for calibration purposes. Ten ‘zero’ and ten ‘flat lamp’ exposures were taken each night. The ‘zero’ exposures were used to remove the bias voltage in the CCD detector so that accurate stellar flux levels could be deduced. A flat lamp frame is a exposure with a effectively flat white light source uniformly illuminating the CCD chip. Flat lamp frames measure the sensitivity variation from pixel to pixel across the CCD chip. Moreover, the echelle grating was moved slightly each night in order to reduce the effect of any flat field artifacts.
Th-Ar comparison frames were taken every two hours throughout each night.
The Th-Ar exposures provide a known comparison spectrum to help determine the wavelength scale in the observed spectra. Dark frames were taken on the first night of each run. Dark frames are CCD images with no light shinning on the chip with the same exposure time as the longest stellar exposure, thereby giving the maximum contribution from thermal noise. They were used to determine the thermal noise level in the stellar exposures so that it could be subtracted if necessary. However, the thermal noise level was found to be very low and did not affect our results. The setup yielded a spectral resolution of about 185,000 or 1.6 km s−1 for observation of the primary sample, as determined from the full widths at half maximum (FWHM) of the thorium lines in the comparison spectra.
2.2.2 ESO Observations
For the stars in the sample too far south to be observable from McDonald Ob- servatory in Texas, data were obtained from the European Southern Observatory
Science Archive Facility. These spectra were acquired by the Ultraviolet and Visual
21 Echelle Spectrograph (UVES) of the Very Large Telescope (VLT) at Cerro Paranal,
Chile. The majority of UVES data (from programs listed in Table 2.3) have S/N
ratios that are comparable to those of our McDonald spectra. These observations
employed central wavelength settings at 4370 A,˚ and 3460 A˚ allowing spectral cov- erage from 3050 to 5000 A.˚ The spectral resolution of the UVES data varies in the different observing programs. This is due to changes in slit width and CCD binning.
The highest resolution data of the target in each program is listed in Table 2.3.
There are three southern sight lines (HD 96670, HD 96715, and HD 164270) that were observed at La Silla but on the New Technology Telescope (NTT) telescope with the ESO Multi-Mode Instrument (EMMI), but with a much lower resolution (Dekker et al. 1986). These sight lines have FUSE data (Table 2.4) but because of the lack of high resolution visible data, they are not discussed further.
2.2.3 FUSE Observations
The UV data for our sample of sight lines were selected from archival FUSE ob- servations. The FUSE instrument operates in the wavelength region 905 to 1187 A.˚
With one exception, all the FUSE observations were taken through the low resolu- tion (LWRS) aperture with an average resolution of 17,000. HD 97253 also has an observation through the medium resolution (MDRS) aperture with a sightly higher resolution of 21,000. Table 2.4 lists the details of the FUSE observations such as the sight line, name of each data set, the number of exposures, and the exposure times.
These UV data from FUSE are the primary source for our measurements of N (H2) and the corresponding excitation temperatures, which are obtained from the spectral synthesis of the (2-0), (3-0), and (4-0) bands of the Lyman B-X transition of H2 (see Federman et al. 2005 and Sheffer et al. 2008). The UV data from FUSE are also our primary source for measurements of CO as the observations span many absorption bands from CO (Sheffer et al. 2003; Crenny & Federman 2004). The CO bands used
22 Table 2.3: Observational Data for ESO UVES/VLT Targets
Star Program ID Exposure Optimal Time (s) Resolution HD 92964 076.C-0431(B) 3900 68000 085.C-0799(A) 6270 68000 HD 97253 266.D-5655(A) 1577 74500 HD 108927 71.C-0513(C) 3923 107200 HD 149404 65.I-0526(A) 180 71100 194.C-0833(H) 5580 107200 HD 151932 60.A-9800(D) 2040 107200 266.D-5655(A) 6910 74500 HD 168076 71.C-0513(C) 3872 107200 083.D-0066(C) 2400 42300 084.D-0481(C) 1200 42300 HD 170740 65.I-0526(A) 420 71100 71.C-0367(A) 7800 107200 082.C-0566(A) 2400 77800 194.C-0833(A) 3960 107200 194.C-0833(C) 3000 71100 194.C-0833(E) 5620 107200 194.C-0833(H) 6920 107200
23 in our study are listed in Table 2.5.
A FUSE observation typically consists of several exposures. Each exposure con-
sists of data taken on four independent detector segments: 1A, 1B, 2A, and 2B.
Each of these partially overlapping segments of the FUSE detector has two channels,
lithium fluoride (LiF) and silicon carbide (SiC). This yields eight overlapping spectra.
The different FUSE segments are listed in Table 2.6. There is a large signal-to-noise
variation between spectra from the different sections of the detector. Whenever there
was a segment with significantly lower signal-to-noise, meaning less than half of the
signal-to-noise in the best segment, it was not included in the combining process.
2.2.4 HST Observations
Our sample included six sight lines with archival HST data. In addition to the
CO bands mentioned previously, the HST spectra also included several multiplets of atomic carbon between 1150 and 1700 A.˚ The available A-X bands of CO and C I multiplets were extracted from the HST data. There were three different settings used in the HST observations. HD 108 was the only sight line observed with the
0.1 x 0.03 aperture and E140H grating with the highest resolution, 143,000. HD 108 and HD 46223 were both observed with the 0.2 x 0.2 aperture and E140M grating with the lowest resolution, 38,500. However, the lowest resolution settings encompass additional A-X bands of CO. The four other HST sight lines were observed with the 0.2 x 0.2 aperture and E140H grating with a resolution of 82,500. Additional information about the HST observations such as the names of each data set used, as well as exposure time, and the central wavelength of each observation are listed in
Table 2.7.
24 Table 2.4: Observational Data for FUSE Targets
Star Dataseta Num. Time Star Dataseta Num. Time Exp. (s) Exp. (s) Northern Targets HD 25638 F9270501000 4 8963 HD 108 D0640101000 8 3084 F9270502000 3 8721 Z9010101000 10 4849 F9270503000 6 11598 HD 5689 E0820601000 4 13777 F9270504000 4 11218 E0820602000 6 18182 HD 45314 P1021301000 10 5515 HD 12882 B0970501000 6 15569 HD 46223 C1680302000 17 7115 U1013201000 3 4584 HD 192303 C1710401000 4 3647 HD 13841 Z9010601000 7 4678 HD 192641 C1710201000 5 4499 HD 13969 D9020201000 3 8308 HD 193793 Q3030101000 4 9191 HD 14053 D9020501000 3 9029 Q3030102000 4 5491 BD+ 56◦ 0501 D9020701000 10 22712 HD 197770 U1042302000 4 13183 BD+ 56◦ 0563 B0290101000 6 11402 A1181301000 4 6007 HD 14443 D9021201000 6 20230 M1141401000 1 2258 HD 14476 D9021301000 5 15042 HD 214419 G0250101000 2 6333 U1016001000 4 9464 G0250102000 3 11028 BD+ 56◦ 0578 B0290601000 5 14442 HD 217086 E0820801000 3 4966 U1016101000 3 5945 U1046401000 5 9529 HD 14947 E0820201000 4 5503 Southern Targets HD 15629 E0820401000 2 3139 HD 92964 Z9013101000 5 1824 HD 16691 E8050101000 4 12178 HD 96670b P1024201000 3 4273 HD 17505 B0970301000 7 8906 HD 96715b P1024301000 7 4597 HD 17520 B0970201000 3 6671 HD 97253 H0230301000 5 2734 BD+ 60◦ 0586 B0970401000 5 14710 E8050401000c 17 6113 HD 19243 B0970101000 2 3590 HD 97484b E0200801000 7 11355 HD 19820 F9270401000 4 778 HD 108927 Q1010404000 2 6494 F9270402000 5 2205 U1084401000 5 14298 F9271401000 6 5615 HD 149404 P1161702000 39 17850 F9271402000 8 25473 HD 151300b E8050501000 5 15748 F9271403000 5 3047 E8050502000 7 20231 F9271404000 6 6488 E8050503000 8 13700 HD 25443 M7171001000 4 4906 E8050504000 9 23684 M7171002000 2 2370 HD 151932 P1170801000 5 2284 M7171003000 2 2504 HD 164270b P1171001000 3 4966 Z9011601000 3 4625 HD 168076 P1162201000 3 6601 HD 170740 P2160601000 6 2328 aUnless otherwise noted, LWRS was the aperture used, and the central wavelength was 1046 A.˚ bNo corresponding high resolution visible data. cThe aperture used in this dataset was MDRS.
25 Table 2.5: FUSE CO Bands
CO Band R(0) Wavelength f -value FUSE Segments (A)˚ Coverage E-X (0-0) 1076.03 0.063 2bsic 1asic 1alif C-X (0-0) 1087.87 0.123 2bsic 1asic 2alif B-X (1-0) 1123.57 0.00106 2alif 1blif B-X (0-0) 1150.48 0.0067 2alif 1blif
Table 2.6: FUSE Segments
Detector Segment Wavelength Material Side Part Range (A)˚ SiC 1 A 1004-1091 SiC 1 B 905-993 LiF 1 A 987-1082 LiF 1 B 1094-1188 SiC 2 A 917-1006 SiC 2 B 1016-1104 LiF 2 A 1087-1182 LiF 2 B 979-1075
26 Table 2.7: Observational Data for HST Targets
Star Data Set Exposure Aperture Gratings Central Resolution Time Wavelength (s) (A)˚ HD 108 OBIL05010 700 0.2X0.2 E140M 1425 38500 O5LH01010 997 0.1X0.03 E140H 1271 143000 O5LH01020 997 0.1X0.03 E140H 1271 143000 O5LH01030 997 0.1X0.03 E140H 1271 143000 O5LH01040 997 0.1X0.03 E140H 1271 143000 OBKR2A010 600 0.2X0.2 E140H 1271 82500 HD 13841 OBKR8M010 1200 0.2X0.2 E140H 1271 82500 HD 25443 OBKR7G010 1200 0.2X0.2 E140H 1271 82500 HD 46223 OCB6C0030 2626 0.2X0.2 E140M 1425 38500 OCB6C1030 2626 0.2X0.2 E140M 1425 38500 HD 108927 OBKR7V010 1200 0.2X0.2 E140H 1271 82500 HD 170740 OBKR1N010 600 0.2X0.2 E140H 1271 82500
2.2.5 GOT C+ Observations
The stellar data and observational details for the additional sight lines in the GOT
C+ sample are provided in Table 2.8. As noted above, the data at visible wavelengths
were taken at the McDonald Observatory with the echelle spectrograph (Tull et al.
1995). The settings containing Ca II λ3933, Ca I λ4226, CH+ λ4232, CH λ4300, and
CN λ3874 is denoted as ‘Blue’ and has a resolution of 135,000. The setting with K I
λ7699 is denoted as ‘Red’ in Table 2.8 and has a slightly lower spectral resolution than the rest of the McDonald data, 125,000 versus 135,000. References for spectral types and (B-V )0 are provided in the footnotes of Table 2.8. Visual extinctions were derived by assuming AV = 3.1 E(B-V ), and distances were obtained by spectro- scopic parallax as many stars did not originally have definitive Hipparcos or GAIA measurements and are known to approximately 20% based on average uncertainties in magnitudes (∼0.02). Recently, a reanalysis of GAIA parallax measurements yielded distances with errors less than 20%. The analysis presented here is based on the
27 spectroscopic distances. Section 4.6.3 discusses the impact of the revised parallax
distances. Although HD 35652 and HD 60146 were observed at McDonald, Herschel
Space Observatory measurements were not obtained, nor were data for CO or H I emission. These two sight lines are not discussed further.
The GOT C+ project is described in detail by Langer et al. (2010) and Langer et al. (2014). Additional observations for the GOT C+ project include the archival data from the Herschel Space Observatory, the ATNF Mopra Telescope, and the
VLA Galactic Plane Survey (Stil et al. 2006). The GOT C+ data sets are available as a Herschel User Provided Data Product1. GOT C+ pointings are labeled as
GXXX.X+YY, which are the longitude and latitude rounded to one decimal; the
2 2 actual coordinates are given in Table 2.2. The [C II] P3/2 → P1/2 observations from the HIFI (Pilbratt et al. 2010; de Graauw et al. 2010) instrument on-board the
Herschel Space Observatory have an angular resolution of 12 arcsec. The observations of the J = 0 → 1 transitions of 12CO, 13CO, and C18O from the ATNF Mopra
Telescope have an angular resolution of 33 arcsec. The H I 21 cm observations from the VLA Galactic Plane Survey (VGPS; Still et al. 2006) have an angular resolution of 1 arcmin. The species observed in emission C+, CO, and H I have a velocity resolution of 0.8 km s−1 while C18O has a velocity resolution of 1.6 km s−1.
1http://www.cosmos.esa.int/web/herschel/user-provided-data-products 28 Table 2.8: Stellar Data for GOT C+ Targets
Star l b B V E(B-V ) Distance1 Distance2 Spectral Exp time Exp time (Mag) (Mag) (Mag) (pc) (pc) Type Blue (sec) Red (sec) July 2014 Observations HD 165783 10.45 0.06 8.60a 8.33a 0.46 1500 500 B3/5Ibb 5400 ... HD 165918 10.51 -0.07 8.22c 8.11c 0.26 350 900 B5IV/Vb 4800 ... HD 168607 14.97 -0.94 9.82d 8.28d 1.37 1500 1100 B9Iaepe 7200 ... HD 167498 15.65 0.85 8.07c 7.83c 0.19 300 400 A2IVb 3600 ... HD 167812 15.88 0.59 8.19c 7.97c 0.29 250 250 B9.5Vb 3600 ... HD 169754 20.02 0.24 9.64d 8.60d 1.11 1750 2000 B0.5Iaf 5400 ... December 2012 Observations BD +49◦ 3482 91.54 0.97 9.40c 9.09c 0.22 400 250 A2g 10800 3600 29 BD +49◦ 3484 91.65 1.04 9.70c 9.43c 0.28 450 400 A0g 7200 ... HD 240179 109.91 -0.04 9.50i 9.28i 0.39 950 700 B5Vg 9000 3600 HD 240183 109.93 -0.05 9.97j 9.85j 0.29 850 1050 B5Vg 12300 7200 HD 35652 173.05 -0.03 8.57k 8.39k 0.40 2250 700 B3Vnnel 7200 3600 HD 47073 207.09 -1.05 8.57c 8.34c 0.19 350 200 A5/7h 5400 3600 HD 260737 207.11 -1.06 9.27c 9.18c 0.02 450 350 A2g 16200 5400 HD 55469 225.24 -0.42 9.57c 9.35c 0.24 250 400 A2IIIh 16200 3600 HD 55981 225.36 0.11 10.02m 9.87m 0.23 1950 650 A0IVh 17700 7200 HD 60146 234.31 0.12 9.48c 9.49c ... 1250 ... B8II/IIIb ... 5400 October 2017 Observations HD 174509 32.36 -0.12 9.48g 9.27g 0.32 550 500 B3V(n)h 14000 7200 1 2 a b — Parallax distance from GAIA DR2 using Bailer-Jones et al. (2018). — Spectroscopic distance. — Jaschek & Egret (1982); — Houk & c d e f g Smith-Moore (1988); — Høg et al. (2000a); — Ducati (2002); — Walborn & Fitzpatrick (2000); — Morgan et al. (1955); — Cannon & Pickering h i j k l m (1993); — Houk & Swift (1999); — ESA (1997); — Kislyuk et al. (2000); — Hohle et al. (2010); — Guetter (1968); — Claria (1974) 2.3 Data Reduction
The visible data reduction was performed with Image Reduction and Analysis
Facility (IRAF) software using standard routines (National Optical Astronomy Ob-
servatories 1999). Dark and bias frames were combined and removed from all other
exposures prior to cosmic ray removal and other processing. Flat fields were nor-
malized and divided out of the exposures. A wavelength scale was established using
Th-Ar lamp exposures (resolving power (λ/∆λ) ∼ 185,000) and applied to all other exposures taken at that time. Afterwards, the calibrated exposures were corrected for their Vlsr offset and combined.
2.3.1 McDonald Observatory Data Reduction
The data were reduced with the standard NOAO echelle data reduction packages and routines in IRAF. The reduction procedure is similar to one described by Knauth
(2001), Pan (2002), and Ritchey (2009). There were seven basic steps involved in the data reduction procedure.
The first step was to remove bad pixels and cosmic ray blemishes from the data.
The cosmicrays task was used for this purpose in both stellar and comparison lamp exposures. Stellar exposures underwent 100 passes using a detection threshold of 50 times the median of the surrounding pixels. The shorter comparison lamp exposures only had to undergo 10 passes using a detection threshold of 25 times the median of the surrounding pixels. The median task was used to process flat lamp exposures.
This routine computes the average of a specified number of pixels along both rows and columns of the detector, and replaces any pixel that is higher or lower than the user-defined standard deviation from the average with the mean within the box.
The second step was to correct for bias voltage on the CCD chip. An average bias
(zero) frame was created for each night with the zerocombine task. The average bias
30 frame was subtracted from all raw stellar, comparison, and flat images. An average
dark image was also created for each observing run, with the darkcombine task, and
the average dark image was compared with the average bias frame statistically using
the imstat task. We found that the average dark frame was essentially identical to the
average bias image for all runs because the detector was cooled with liquid nitrogen
to a temperature of 110 K. Therefore, no thermal noise correction was performed on
any data.
The third step was to minimize the effect of pixel to pixel variation in sensitivity
across the CCD detector. First, the task apscatter fits and removes scattered light
within the spectrograph through the use of a low order polynomial. The locations
of the apertures and their widths on the detector were marked with the IRAF task
aptrace. These tasks were performed on all stellar exposures and on the combined
flat field image. Then, the average flat frame was normalized to unity with the
apnormalize task. This normalized flat frame was then divided into each stellar
image for further processing, minimizing the variation in pixel to pixel sensitivity
across the CCD detector.
The next step was to extract one dimensional (1D) spectra from the two dimen-
sional images. The pixels perpendicular to the dispersion axis in each order of the
echelle spectra were summed together with the task apsum for each stellar and
comparison exposure. This yielded 1D extracted spectra to which an appropriate
wavelength scale was applied.
The next step created a template for wavelength calibration. Emission lines in
the extracted 1D comparison spectra were identified and a wavelength solution was
created for each comparison spectrum using the task ecident or ecreident. This process requires several (ideally at least five) emission lines in each order (typically eight orders) to be identified and marked for the task to work accurately. The remain- ing unmarked emission lines in the spectrum are automatically added through the use
31 of a Th-Ar line list. Once these identifications were made, a dispersion solution was
calculated and an accurate wavelength scale was applied to the data with the tasks
refspec and dispcor.
The next step used the dopcor task and corrected for the Doppler motion arising from the Earth’s rotation and orbital motion around the Sun, yielding a scale in the velocity of the Local Standard of Rest (Vlsr= c(λ-λlab)/λlab), where λ is the observed wavelength and λlab is the laboratory wavelength of the transition and c is the speed of light.
The final step created a summed spectrum for each interstellar feature of interest, where the continuum was normalized. Spectra with interstellar absorption lines and sufficient continua on both sides were created with the task scopy for each stellar exposure. Each spectrum was carefully examined for artifacts and blemishes that re- mained before the summation. The cosmic ray blemishes were subsequently removed unless the blemishes coincided with interstellar lines in the spectrum. Then, the spec- tra for the same species toward the same star were co-added together with the task scombine to obtain a final spectrum with higher signal-to-noise (S/N). The co-added spectra were normalized, by fitting low order Legendre polynomials to regions free of interstellar absorption.
32 2.3.2 ESO Data Reduction
All ESO data sets were reduced with the UVES pipeline software by Dr. Adam
Ritchey. The optimal extraction mode was used in all cases except when the observa- tions employed an image slicer in which case the average extraction method was used.
Master bias and master flat frames were created each night. Th-Ar lamp spectra were acquired alongside the science data and were used for wavelength calibration. After extraction, all echelle orders from a given exposure were combined and the resulting spectra were shifted to the LSR frame of reference. Finally, all individual exposures of the same target obtained with the same instrumental setup were co-added. The co-added UVES spectra were normalized to the continuum in much the same way as were the McDonald data.
2.3.3 Known FUSE Detector Effects
The flux from individual segments, described in Section 2.2.3, must be combined together to get the final spectrum. However, there are several effects that have to be taken into consideration before creating a combined spectrum.
The apparent movement or shifting of the target can create differences between each exposure. This can come from the target moving around within the aperture between exposures. Any jitter from the movement of the spacecraft can also add to the movement of the target. These effects are largely corrected for in the reduction pipeline, but an offset in the wavelength scale for each exposure can remain (Section
7.6 of the FUSE Data Handbook, Dixon et al. 2007). Offsets on the order of 0.1 A˚ between exposures have been observed in some of our data.
There are several flux effects that also create problems when combining exposures.
Thermal effects cause a slight flexing of the optical bench on orbital timescales, re- sulting in the drifting of the relative alignment of the four FUSE channels by several
33 arc seconds with respect to each other (FUSE Data Handbook). This is referred to as Channel Drift, where the source can drift in and out of the channels during an exposure, resulting in lost counts. Observations in the MDRS and high resolution
(HIRS) apertures often lost flux due to this problem. LWRS observations were less affected by flux loss, but could show offsets in wavelength scale between exposures
(Section 7.5.5 of the FUSE Data Handbook).
A more subtle flux effect comes from grid wires in front of the detectors that occa- sionally cast shadows in detected spectra, artificially lowering the detected flux level over local regions. These shadows were referred to as worms. This effect is important to remove because of the possibility for the shadow showing up as a artificial compo- nent contaminating the line of interest. Luckily, their presence can be investigated by comparing data for overlapping wavelength ranges in different detectors. Section 4.3.4 of the FUSE Instrument Handbook and Section 7.3.2 of the FUSE Data Handbook provide further details.
Scattered light can also affect the local background levels on the FUSE detectors and can be responsible for producing a non-zero “background,” even in optically thick absorption lines. This primarily impacts SiC data where solar-scattered light is more severe (Section 7.1.3 of the FUSE Data Handbook). Fortunately, there are several indicators for detecting and identifying these anomalies. The 2-D images of the detector typically show evidence of scattered light or other unwanted signals. In addition, “trailer files” of the calibrated data are included and show indications of anomalies during processing of a data set.
34 2.3.4 FUSE Data Reduction
As part of the FUSE project close-out, all FUSE data were re-processed using the final version of the CalFUSE calibration pipeline software package, CalFUSE
3.2. The pipeline creates calibrated spectra for each of the channels in each of the segments for each exposure. The raw and calibrated FUSE data files and ancillary
files are archived with MAST, originally the acronym for the Multimission Archive at the Space Telescope Science Institute but in 2012 was renamed the Barbara A.
Mikulski Archive for Space Telescopes1. All further processing of FUSE observations used this final reduction pipeline as the starting point.
First, the observations were rebinned by a factor of four to improve signal-to- noise. The next step is to combine spectra from the same segments in a observing run with scombine, with any extremely low S/N spectra removed before proceeding.
Unfortunately, there are variations in the observed wavelength offset between the different sections of the detector. In some cases there are further detector effects that must be removed before the data can be analyzed, such as worms and detector grid shadows, as noted in Section 2.3.3. A full catalog of factors potentially impacting
FUSE data quality are given in Chapter 7 of the FUSE Data Handbook.
To combine the segments together, we must take these offsets into account. The poff task determines the shifts (pixel offsets) between a given spectrum and a list of other spectra. This information is used by the specalign task to combine the individual spectra that are shifted relative to one another to create an appropriately combined flux spectrum.
Next, the spectra from multiple exposures are compared, with any low S/N spectra removed before proceeding. Each of the spectra have already been shifted, but that does not mean they have been shifted in identical ways. They need to be aligned and
1http://archive.stsci.edu/fuse/
35 combined again in the same manner as before, poff followed by specalign. If there are multiple observations, the next step is to combine them into one spectrum using the same steps outlined previously and removing any spectra with low S/N. The final spectrum can now be normalized by fitting a low order Legendre polynomial to the continuum as part of the fitting process, as described in the next chapter.
2.3.5 HST Data Reduction
Hubble archival data are processed using an On-The-Fly Reprocessing (OTFR) system. The OTFR system reconstructs the data files with updated headers and calibrates data when processing a user’s request for the data from the archive. This is the starting point we used for all HST data. Next, all the overlapping orders in each exposure from the archive are combined into a single spectrum with the splice task. Then the spectra are converted from a FITS 3-D binary table into the multispec format image for IRAF with the tomultispec task. When there are multiple exposures, their spectra can be combined by summing them together with the scombine task.
2.3.6 GOT C+ Reduction
The GOT C+ project is described in detail by Langer et al. (2010), Langer et al.
(2014), and Velusamy et al. (2010). The details of reduction of the emission data are described in Pineda et al. (2013). The McDonald Observatory spectra, taken to accompany GOT C+ pointings, were reduced as previously described in Section 2.3.1.
36 Chapter 3
Component Fitting
After the spectra were reduced, each feature was isolated along with a few A˚ of the spectrum on each side the feature. This allowed the continuum around each feature to be fit with a low order Legendre polynomial for normalization, typically with an order of less than four. This was performed within IRAF through the t sub-task of the splot program to interactively fit a function to the spectrum using the ICFIT mode.
Any points that were not associated with the feature and had significant residual differences with the fit (>1.5 sigma) were rejected. This interactive selection/rejection was used to exclude spectral defects and lines from the fitting process. Dividing the spectrum by the smooth fitted continuum resulted in a normalized spectrum.
In most of the visible spectra, the regions surrounding the features were free from stellar absorption features and contamination from other interstellar species. Some of the CO bands did have absorption lines from other interstellar species and stellar contamination that had to be removed also before a normalized continuum level could be established. In some cases, the continuum level had to be iteratively refit to remove stellar contamination or account for absorption from the lines of other species such as H2. After normalization, the features in each spectrum were fit with a heavily modified version of the ISMOD fitting routine (Sheffer, unpublished), discussed in Section 3.1.
37 The lines of species seen at visible wavelengths were always fit first and used as the starting point for fitting the lines of species found in the lower resolution UV data.
Section 3.2 describes the first step in the fitting process, fitting lines of the species at visible wavelengths. Sections 3.3 and 3.4 discuss fitting H2 and CO, respectively. The last section of the chapter, Section 3.5, discusses fitting the lines of other species such as neutral carbon and species in emission from the GOT C+ data, which will provide additional information for my analysis.
3.1 ISMOD
ISMOD is a root mean square (rms) error minimization fitting routine written in FORTRAN, originally created by Y. Sheffer. A synthesized absorption profile of each component is created by using the velocities in V lsr, Doppler b-values, relative fractions (f ) and total column density (N ) as potential free parameters during the
fitting process. For species with low lying rotational or fine-structure levels, each level is treated separately. In an effort to avoid over-fitting features, we try to use the fewest components needed for the residuals of the fit minus data to reach the same rms as the background continuum.
The ISMOD routine assumes a Voigt profile function for each absorption com- ponent. It convolves the intrinsic Voigt profile of each component with an Gaussian instrumental profile, where the width of the Gaussian is determined by the resolving power of the instrument (see Black & van Dishoeck 1988). Each iteration varies one of the free parameters slightly and checks if the overall residual rms between the data and the fit has increased or decreased. These iterations continue until the change in parameters for each step is less than one part in 105.
There are several potential issues with this method of fitting. First, by using rms minimization, there is a danger of finding a local minimum in the available
38 parameter space. Second, the larger the parameter space, the greater the possibility of degenerate solutions. Both of these issues can be mitigated by choosing a better starting point based on some prior knowledge about the rough solution or by limiting the available parameter space for the fitting.
There is also the possibility of the best fit parameters corresponding to a nonphys- ical solution. After each run of the program, the parameters from its rms minimized solution and the resulting fit are evaluated to verify that they are physically reason- able. The details of what is physically reasonable for each species are discussed in the corresponding sections. If a solution is found to be nonphysical or contradicts other information, then the starting point or the allowed limits would be modified and the
ISMOD program would be restarted. These iterations continue until the best version of the fit is found.
3.2 Fitting Ground-based Spectra
The lines from species in the visible were always fit first and used as the starting point of the first iteration in fitting the lines of species found in the HST data and the lower resolution FUSE data. The wavelengths and f -values for the lines in the visible are given in Table 3.1. Initially, the lines of species with the simplest structure was fit first. Simpler structure is usually found in the probes of denser gas. This profile was then used to help fit the more complicated profiles of species probing more diffuse gas. A helpful schematic of this idea is found in Figure 6 of Pan et al.
(2005), which shows overlapping layered locations of tracers at visible wavelengths within a cloud. When detected, the CN profile is usually much less complicated than profiles of other species for a given sight line, as CN absorption traces relatively dense
−3 gas (nH ∼ 300 cm ). This is reflected in the allowed range of b-values for CN, 0.3 to 1.5 km s−1 (Pan et al. 2004). In the regions of the cloud surrounding this gas, we tend
39 to find absorption from CH. The column densities of CN components are correlated
with column densities of corresponding CH components surrounding the denser CN
regions. We also tend to find CH components associated with lower density CH+ components. This is because CH has two production routes, one that dominates at higher density and one that dominates at lower density (Pan et al. 2005). This results in the component structure of the CH profile slightly more complicated than the structure in the CN profile. The distribution of CH also results in larger allowed b-values, 0.5 to 2.5 km s−1.
When observed, the K I profile provides the next step in fitting the complicated structure for some of the species found in more diffuse gas. Welty & Hobbs (2001) found an essentially linear relationship between N (K I) and N (CH) by using total column densities along lines of sight. The column densities of N (K I) and diffuse
N (Ca I) are also shown to be well correlated by Welty et al. (2003).
The most diffuse gas is probed by Ca II, which tends to have the most complicated structure of the species discussed here. Ca II profiles are typically deeper and reside in gas more widely distributed in velocity than the other atomic species considered here (Pan 2002; Pan et al. 2004, 2005; Sheffer et al. 2008; Ritchey 2009). The allowed b-values for Ca II are the largest of the species discussed here, 1.0 to 3.5 km s−1. The
line widths of individual Ca II components track those of the dominant ion in diffuse
clouds (Welty et al. 1996; Ritchey 2009).
Each of the resulting fits were crosschecked with fits for all the lines from other
available species (Ca II, Ca I, CH, CH+, CN, and K I) from the same line of sight. If
there was disagreement, all the spectral features were iteratively refit until an overall
agreement was reached. Figures 3-1 and 3-2 illustrate the agreement and differences
seen in the spectra of all the observed species in an individual sight line. Above the
spectrum for each species, red tick marks are shown to indicate the locations of the
observed velocity components. Figure 3-1, toward HD 13841, reveals a component in
40 CH and CH+ near -40 km s−1 from the Perseus Spiral Arm of the Milky Way. The
Ca II profile toward HD 13841 also shows components extending into the Perseus Arm.
Figure 3-2 shows a much less extended distribution of Ca II components toward HD
197770. All other sight lines not used as examples have corresponding figures showing their stacked spectra of observed species in Appendix A.
The results for component column densities, structure, and b-values from the profile fitting are given over several tables. Table 3.2 lists the results from McDonald
Observatory for CH, CH+, and CN. When CN is not detected, a three sigma upper limit for the total column density of CN is listed at the Vlsr of the strongest CH component. If J = 1 is not detected in CN, it is assumed to be 50% of J = 0. Table
3.3 lists the component results from the McDonald Observatory for fits to the spectra of Ca II. Table 3.4 presents the component and total equivalent widths. Uncertainties in the total equivalent width are determined from the full with at half max and the rms uncertainty. Uncertainties in the total column density are determined from the total column density and the ratio of the total equivalent width uncertainty and the total equivalent width. The additional tables list the results of the species observed in the visible for each additional data set, Table 3.5 shows the results for the CH,
CH+, and CN for the sight lines in the Southern Hemisphere obtained by Dr. Adam
Ritchey and Table 3.6 lists the results for the GOT C+ sight lines.
41 Table 3.1: Spectroscopic Data at Visible Wavelengths
Species Band λ(A)˚ Line f -value CN B-X (0,0) 3874.602 R(0) 0.0342 CN 3874.000 R(1) 0.0228 CN 3875.760 P(1) 0.0114 Ca II 3933.663 0.6267 Ca I 4226.728 1.77 CH+ A-X (0,0) 4232.548 R(0) 0.00545 CH A-X (0,0) 4300.313 R(0) 0.00510 K I 7698.965 0.3327
42 Figure 3-1: The figure shows the stacked spectra of the species seen toward HD 13841. The species for each section of the spectrum is indi- cated by a label in the bottom left corner. The fit is indicated by solid lines while the data appear as small ‘x’s. The Vlsr of the components for each species are shown as red tick marks above each region of the spectrum. A dashed blue line at the contin- uum level is provided for comparison. A fitted line is not shown for the CN spectrum whenever only upper limits are reliably de- termined.
43 Figure 3-2: The figure shows the stacked spectra of the species seen toward HD 197770. CN is detected and a fitted line is shown for the R(0) line of the CN spectrum. The R(1) line of CN can also be seen as a small feature to the left of the R(0) branch. Figure 3-3 shows the complete fit to all the CN lines.
44 Figure 3-3: The figure shows the spectrum of CN toward HD 197770. The R(0), R(1), and P(1) lines of CN are detected and a fitted solid line is overlaid over the CN spectrum.
45 Table 3.2: CH, CH+, and CN Results from McDonald Observatory Observa- tions
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 108 -15.4 3.1(0.7) 2.2 -16.4 2.1(0.7) 1.9 ...... -9.6 2.6(0.4) 0.5 -11.7 2.0(0.5) 1.0 ...... -4.8 4.1(0.6) 2.1 ...... -5.1 1.5(0.2) 1.3 0.8 1.7(0.7) 2.5 -0.8 3.9(0.9) 3.2 ...... 1.9 2.5(1.2) 2.8 ...... 5.2 1.4(0.7) 2.5 ...... 6.9 1.1(0.2) 0.8
46 Total 12.9(3.7) 10.4(3.4) 2.6(0.6)
HD 5689 -20.3 2.6(1.3) 0.5 -20.2 3.4(1.8) 1.0 ...... -12.6 9.3(2.2) 1.8 -13.2 10.6(2.6) 1.7 -12.6 ≤4.4 ... -8.8 3.8(1.2) 0.6 -8.0 12.1(3.2) 2.1 ...... -1.4 4.3(2.7) 2.5 -0.4 5.4(1.9) 1.0 ...... Total 17.4(7.1) 31.5(9.8)
HD 12882 -5.4 4.5(0.7) 0.7 ...... -2.4 4.1(0.8) 1.2 -1.9 2.1(0.9) 1.4 ...... 0.2 7.6(1.6) 2.3 ...... 0.2 ≤2.7 ... 2.1 4.6(0.7) 1.3 2.4 1.9(1.0) 1.8 ...... 5.0 3.2(1.4) 2.0 ...... Total 20.9(3.8) 7.2(3.3) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 13841 -41.6 4.8(0.4) 1.1 -42.6 2.5(0.4) 1.5 -41.6 ≤0.7 ...... -16.0 2.6(0.8) 3.4 ...... -13.5 0.7(0.3) 1.4 ...... -7.9 3.6(0.6) 2.6 ...... -4.8 1.6(0.4) 1.4 ...... -2.8 1.9(0.6) 2.5 ...... -0.7 3.0(0.7) 2.8 ......
47 Total 6.6(1.3) 14.0(3.4)
HD 13969 -15.0 2.2(0.7) 0.6 -15.4 5.4(1.4) 2.1 ...... -10.8 4.7(0.9) 1.3 ...... -10.8 ≤1.6 ...... -8.0 5.2(1.7) 2.8 ...... -1.7 5.3(1.8) 2.6 ...... Total 6.8(1.7) 15.9(4.9)
BD+ 56◦ 0501 -16.5 3.3(2.5) 2.5 ...... -9.2 5.0(1.2) 0.5 ...... -9.2 ≤2.9 ... -4.9 2.1(1.3) 0.6 -4.7 7.6(2.9) 2.5 ...... 0.9 5.4(1.7) 1.2 ...... 16.0 4.2(1.6) 1.0 ...... Total 10.4(5.8) 17.2(6.1) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) BD+ 56◦ 0563 -12.3 1.7(0.9) 0.5 ...... -8.4 3.6(2.3) 2.3 -8.0 6.7(1.9) 1.8 -8.4 ≤1.5 ... -5.5 3.5(2.0) 2.4 -5.5 4.4(3.7) 2.1 ...... -3.5 4.1(1.5) 1.3 ...... -1.3 2.1(1.2) 0.5 ...... Total 10.9(6.5) 15.2(7.1) HD 14053 ...... -13.8 5.3(1.2) 1.0 ...... 48 -8.6 9.4(1.4) 2.2 -9.2 8.2(2.2) 2.4 -8.6 ≤2.3 ...... -6.1 7.7(1.2) 1.0 ...... -3.7 2.3(0.7) 0.5 -3.5 4.3(1.2) 1.4 ...... 0.0 4.1(0.7) 0.5 -0.1 7.5(1.5) 1.6 ...... Total 15.8(2.9) 32.9(7.5) HD 14443 ...... -10.1 5.2(1.1) 1.5 ...... -6.2 3.4(1.4) 2.3 -5.3 6.2(1.4) 2.1 -6.2 ≤1.2 ... -1.0 2.2(0.8) 0.5 0.0 7.7(1.3) 1.8 ...... Total 5.6(2.2) 19.0(3.7) HD 14476 -7.9 3.3(1.7) 0.5 ...... -4.3 3.2(1.3) 0.5 -6.0 7.0(4.1) 2.5 ...... -0.6 3.7(2.9) 0.6 -0.4 6.8(4.2) 2.5 -0.6 ≤2.2 ... Total 10.1(5.8) 13.7(8.3) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) BD+ 56◦ 0578 -23.2 2.5(1.5) 0.5 ...... -12.4 5.1(1.9) 1.2 ...... -12.4 ≤3.2 ... -6.6 2.1(1.2) 0.5 -7.1 4.1(3.0) 3.5 ...... -4.2 3.3(2.2) 1.0 ...... Total 9.8(4.5) 7.4(7.7)
HD 14947 -11.0 9.9(1.2) 2.2 -11.5 8.7(1.4) 2.5 ...... -7.7 11.0(0.9) 1.3 -7.3 5.8(1.4) 2.0 -7.7 ≤1.7 ... 49 -4.8 4.5(0.8) 2.1 -3.9 5.9(1.1) 1.6 ...... Total 25.4(3.3) 20.4(3.9)
HD 15629 -5.1 5.5(1.6) 0.9 -5.9 4.5(2.3) 1.0 -5.1 ≤3.6 ...... -2.5 4.8(2.7) 1.3 ...... -0.3 4.8(1.4) 0.5 ...... 6.7 4.1(2.3) 1.1 ...... Total 10.3(3.1) 13.5(7.2)
HD 16691 -46.9 8.2(2.3) 1.6 ...... -12.7 9.5(3.1) 2.5 -12.1 8.1(3.1) 1.2 -12.7 ≤8.3 ...... -4.1 22.4(7.1) 3.5 ...... 4.4 4.0(2.3) 1.5 4.9 14.1(3.8) 1.6 ...... Total 21.6(8.0) 44.6(14.1) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 17505 -6.8 6.0(0.6) 1.3 -5.4 26.5(1.3) 2.1 ...... -4.3 7.8(0.5) 1.0 -3.4 5.0(1.1) 3.5 -4.3 ≤1.4 ... -1.9 4.7(0.4) 0.5 ...... 1.1 2.4(0.4) 0.5 1.3 3.4(1.4) 2.6 ...... 5.4 1.6(0.4) 0.5 ...... Total 22.5(2.6) 34.9(6.0) 50
HD 17520 -5.9 4.5(0.9) 0.8 ...... -2.5 4.7(0.9) 1.2 -1.8 6.4(1.2) 1.0 -2.5 ≤2.2 ...... 1.6 3.5(1.2) 1.1 ...... Total 9.1(1.8) 9.8(2.5)
BD+ 60◦ 0586 -6.4 2.7(0.6) 0.5 -6.1 5.5(2.2) 3.5 ...... -3.6 1.3(0.6) 1.0 ...... -2.2 6.4(1.2) 1.8 -2.2 1.4(1.3) 1.1 -2.2 ≤1.7 ... 0.0 5.5(1.2) 2.5 0.3 3.6(1.1) 1.3 ...... 3.1 5.0(2.0) 1.7 ...... Total 14.6(3.0) 16.8(6.9) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 19243 ...... -11.9 2.0(0.7) 2.2 ...... -10.0 4.1(1.0) 2.2 -9.5 4.7(0.7) 1.2 ...... -7.7 3.6(0.6) 2.2 -6.3 2.1(0.6) 1.0 ...... 0.5 4.9(0.8) 2.3 0.0 1.2(0.9) 3.3 0.5 ≤0.8 ...... 2.6 8.5(1.5) 3.5 ...... Total 12.6(2.3) 18.5(5.8)
51 HD 19820 ...... -19.1 4.4(2.0) 3.5 ...... -14.9 17.2(1.1) 1.7 -15.0 5.8(1.5) 2.2 -14.9 ≤2.46 ...... -5.8 2.8(0.8) 1.0 ...... -3.4 2.9(0.9) 1.5 -2.8 4.0(0.9) 1.0 ...... -0.8 4.4(0.6) 0.5 0.8 4.3(1.2) 1.7 ...... 2.6 2.8(0.8) 1.1 ...... Total 27.3(4.3) 21.5(6.7)
HD 25443 ...... -14.2 8.7(1.2) 1.9 ...... -10.9 4.9(1.2) 1.6 -10.9 6.7(1.3) 1.6 ...... -8.1 4.6(0.6) 0.8 -7.5 7.1(1.1) 1.4 -8.1 ≤1.6 ... 1.8 3.7(0.8) 0.9 ...... 6.4 3.8(0.7) 0.5 ...... Total 16.9(3.1) 22.4(3.7) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 25638 -16.1 4.1(0.5) 0.8 -16.1 5.9(3.0) 2.7 ...... -10.2 13.6(1.0) 2.3 -10.0 10.6(1.4) 2.5 ...... 2.2 10.5(0.8) 1.8 ...... 2.4 3.8(0.3) 0.6 6.5 11.1(0.6) 1.0 5.3 5.1(1.3) 2.3 5.7 2.0(0.3) 0.5 Total 39.2(3.1) 21.6(5.6) 5.8(0.8) 52
HD 45314 ...... -5.1 1.7(0.8) 1.0 ...... 1.7 6.1(1.6) 2.4 ...... 1.7 ≤1.8 ... 5.1 5.1(0.9) 1.4 ...... 8.3 2.1(1.3) 2.1 ...... Total 11.2(2.4) 3.8(2.2)
HD 46223 1.4 2.4(0.5) 0.6 0.2 2.9(1.5) 2.4 ...... 8.2 7.9(1.2) 2.5 8.1 6.0(1.9) 3.5 ...... 11.3 9.7(0.8) 1.8 10.8 11.7(1.1) 2.2 11.3 ≤1.3 ... Total 20.0(2.7) 20.6(5.2) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 192303 3.4 2.5(1.3) 1.0 4.1 14.2(2.5) 1.9 ...... 8.9 5.6(2.3) 2.5 ...... 8.9 ≤3.8 ... Total 8.0(3.7) 14.2(2.6)
HD 192641 0.0 2.7(1.0) 1.3 ...... 4.0 9.6(1.2) 1.5 3.9 21.8(2.1) 2.4 4.0 ≤1.5 ...... 9.4 5.8(1.6) 1.9 ......
53 Total 12.3(2.6) 27.7(4.4)
HD 193793 -3.5 2.7(0.7) 2.2 -3.2 5.5(1.2) 3.0 ...... 6.2 7.3(0.7) 1.9 5.6 19.4(1.2) 2.6 6.2 ≤0.7 ... 10.1 6.0(0.7) 2.4 10.1 32.3(1.4) 2.8 ...... 13.9 5.9(0.6) 2.1 ...... 18.2 0.8(0.3) 0.5 17.2 4.8(0.7) 1.5 ...... Total 22.7(3.1) 62.0(5.8)
(J =0) / (J =1) HD 197770 ...... -3.5 0.9(0.4) 1.0 ...... -1.1 31.5(0.3) 1.1 -0.7 1.8(0.5) 1.5 -1.4 5.0 (0.2) / 2.0 (0.6) 0.8 2.5 5.6(0.9) 3.0 ...... Total 37.1(2.5) 2.7(0.8) 7.0(0.3) Table 3.2: —Continued
CH CH+ CN a Star Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 214419 -18.8 8.6(0.7) 0.5 ...... -19.5 5.6(0.9) 0.2 ...... -7.9 2.3(0.6) 1.0 ...... -5.6 4.9(1.1) 1.8 ...... -3.9 12.3(1.1) 1.5 ...... -0.4 3.0(1.0) 1.6 -0.9 1.0(0.9) 1.0 ...... 6.9 9.4(1.5) 2.7 6.3 2.2(1.3) 1.5 ...... 54 Total 26.0(4.7) 17.8(3.4) 5.6(0.9)
HD 217086 -21.6 3.6(0.7) 0.5 -20.2 4.1(1.3) 2.1 ...... -15.7 16.1(1.2) 1.7 -14.1 4.4(0.8) 1.1 -15.7 ≤2.5 ... -11.7 9.7(1.0) 1.6 -11.4 7.5(1.2) 1.7 ...... -6.7 7.7(1.1) 1.5 ...... -4.5 4.1(1.4) 2.5 ...... 0.1 10.6(1.1) 1.6 ...... 6.5 2.4(0.7) 0.5 6.4 5.7(1.3) 2.0 ...... Total 36.0(6.3) 40.0(7.2) aUpper limits for CN (J = 0 and J = 1) combined column densities are computed as 1.5 X upper
limit of CN (J = 0) column density at same the Vlsr as the largest CH component. Table 3.3: Ca II Results from McDonald Observatory Observations
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) HD 108 -54.0 0.10 2.5 HD 12882 -12.1 0.14 2.5 -48.5 0.09 2.5 -7.5 0.20 1.4 -42.7 0.16 2.5 -4.7 0.68 1.7 -36.9 0.25 2.4 -2.1 0.62 1.3 -32.2 0.28 2.3 0.0 0.43 1.4 -27.9 0.17 2.3 2.6 0.42 1.3 -23.2 0.11 2.4 4.6 0.30 1.4 -20.4 0.11 1.5 8.7 0.27 1.6 -18.4 0.28 1.4 12.7 0.25 1.6 -15.0 0.55 1.7 16.6 0.09 1.5 -11.9 0.47 1.9 25.1 0.12 2.0 -9.4 0.56 2.0 -5.8 0.75 1.6 HD 13841 -72.0 0.05 2.5 -3.8 0.17 1.4 -67.5 0.23 1.6 -2.0 0.93 2.1 -64.0 0.15 2.5 -0.3 0.44 1.4 -59.1 0.39 1.7 1.5 1.17 1.5 -55.0 0.33 2.5 3.5 0.61 1.4 -52.3 0.46 2.1 5.2 0.51 1.5 -47.7 1.42 2.5 7.7 0.12 1.4 -45.0 0.17 2.5 10.0 0.08 2.0 -42.2 0.69 1.9 15.7 0.16 1.4 -38.4 0.39 2.3 17.6 0.17 2.0 -33.8 0.44 2.1 21.5 0.09 2.0 -30.8 0.24 1.4 23.0 0.02 2.0 -28.0 0.43 2.3 -23.8 0.15 1.7 HD 5689 -58.0 0.15 2.5 -21.0 0.13 2.5 -54.0 0.25 2.5 -16.8 1.06 2.0 -38.8 1.20 2.5 -13.0 0.73 2.0 -33.1 0.41 2.0 -10.0 0.73 1.6 -28.0 0.15 2.0 -7.2 0.94 1.4 -18.7 0.36 1.4 -4.8 1.29 1.9 -14.9 0.26 1.5 -2.8 0.69 1.6 -12.9 2.90 1.5 -0.5 0.40 1.4 -8.5 9.00 1.4 2.0 0.64 1.4 -5.1 1.50 1.8 5.3 0.22 1.4 -1.0 1.71 2.0 8.3 0.09 2.5 1.7 1.12 1.7 17.0 0.03 2.5 5.5 0.31 2.0
55 Table 3.3: —Continued
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) HD 13969 -62.0 0.10 1.4 HD 14053 -61.4 0.28 2.5 -57.0 0.25 2.0 -54.8 0.64 2.5 -54.5 0.17 2.0 -51.7 0.67 2.5 -51.1 0.34 1.4 -45.9 0.74 2.5 -46.0 0.93 2.5 -41.5 0.63 2.5 -43.0 0.76 1.8 -38.3 0.58 2.3 -39.0 0.54 2.4 -34.4 0.61 1.4 -32.9 0.78 2.4 -30.6 0.74 1.8 -27.3 0.06 2.5 -26.0 0.35 2.5 -25.5 0.38 2.5 -21.4 0.40 2.4 -19.6 0.31 2.5 -18.8 0.37 1.5 -14.5 1.08 2.0 -16.0 0.51 1.5 -11.0 1.14 1.8 -13.2 0.54 2.4 -8.0 0.36 1.4 -11.4 0.99 2.0 -6.0 2.69 2.3 -9.2 0.26 1.6 -1.3 1.13 2.0 -7.5 1.29 1.6 0.5 0.33 1.5 -5.7 0.14 1.9 2.9 0.17 1.4 -4.0 1.35 2.0 5.3 0.14 1.5 -2.1 1.12 1.5 -0.3 0.31 1.5 2.9 0.62 2.5
56 Table 3.3: —Continued
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) BD+56◦ 0501 -63.8 0.18 1.8 BD+56◦ 0563 -65.3 0.34 2.5 -57.5 0.27 1.9 -62.8 0.44 2.5 -53.0 0.40 2.1 -54.3 0.27 2.3 -49.2 0.23 1.0 -49.2 0.36 2.4 -44.1 0.64 2.5 -42.1 0.26 1.9 -38.3 0.53 2.5 -38.7 0.35 1.8 -32.5 0.83 2.5 -35.0 0.26 1.6 -28.6 0.50 1.0 -32.1 0.23 1.0 -26.4 0.21 2.0 -28.9 0.27 1.5 -22.0 0.19 2.5 -23.8 0.28 1.9 -17.9 0.33 1.0 -19.2 0.17 1.0 -14.7 0.54 2.5 -16.1 0.23 1.1 -11.6 0.68 1.5 -13.2 0.35 1.2 -8.8 1.10 1.5 -11.7 0.64 2.0 -5.6 1.26 1.1 -9.0 1.16 1.3 -2.1 2.94 1.1 -6.4 0.84 1.1 1.8 0.69 1.9 -3.2 3.55 1.1 6.2 0.32 1.2 -1.0 0.20 2.0 1.5 0.78 1.8 3.8 0.26 2.5
57 Table 3.3: —Continued
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) HD 14443 -70.0 0.10 2.5 HD 14476 -63.0 0.20 2.0 -63.5 0.30 2.5 -59.0 0.26 2.0 -58.5 0.24 2.5 -49.8 0.80 2.5 -53.0 0.40 2.0 -44.5 0.50 2.5 -49.1 0.33 2.5 -38.6 0.55 2.5 -42.7 0.39 2.5 -33.1 0.50 2.5 -38.1 0.56 2.3 -28.5 0.50 2.2 -32.8 0.47 2.5 -23.8 0.50 1.2 -27.2 0.23 1.6 -19.7 0.55 2.5 -24.0 0.20 1.4 -15.4 0.81 1.1 -20.0 0.43 2.4 -11.8 1.03 1.1 -15.7 0.24 1.7 -10.1 1.00 2.0 -12.3 0.98 2.4 -8.1 1.12 2.3 -10.5 0.82 2.3 -6.4 1.30 2.3 -7.3 1.40 2.0 -4.6 1.38 1.8 -5.5 0.01 1.5 -1.0 0.74 1.5 -3.7 1.45 1.7 1.9 0.55 1.1 -0.5 0.34 1.4 3.7 0.30 2.2 1.7 0.53 1.6 5.5 0.20 1.1 3.5 0.17 1.4 5.9 0.18 2.4
58 Table 3.3: —Continued
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) BD+56◦ 0578 -64.7 0.24 2.4 HD 15629 -36.6 0.53 2.5 -59.2 0.24 1.1 -28.8 0.66 2.5 -52.0 0.54 2.1 -22.8 0.50 1.8 -49.5 0.26 1.4 -20.1 0.69 1.5 -41.5 0.28 2.5 -15.2 1.15 1.7 -38.6 0.31 2.1 -10.6 1.00 2.0 -34.5 0.20 1.4 -6.8 0.01 1.0 -30.2 0.18 1.8 -5.3 3.33 2.0 -23.2 0.01 1.1 -2.0 374 1.6 -21.2 0.16 1.6 0.0 0.01 1.3 -16.6 0.28 2.0 1.6 0.01 1.0 -13.1 0.52 1.3 4.3 0.59 1.9 -10.2 2.04 1.0 -6.8 1.51 1.0 HD 16691 -53.4 0.85 2.5 -3.2 3.16 1.1 -47.3 1.47 1.0 -1.7 0.58 1.0 -43.7 1.47 1.5 2.2 0.91 2.2 -40.1 0.72 1.5 -36.6 1.13 1.0 HD 14947 -58.6 0.75 2.3 -32.6 0.95 2.0 -52.5 0.19 2.5 -27.1 0.36 1.0 -44.8 0.25 2.3 -23.1 0.37 1.0 -39.9 0.20 1.9 -18.4 0.32 1.0 -34.8 0.57 2.5 -15.3 0.88 1.0 -29.0 0.34 2.5 -12.1 2.81 2.0 -23.7 0.41 2.5 -7.3 1.25 1.5 -19.8 0.46 2.4 -3.9 1.91 2.0 -15.7 0.59 2.3 -2.0 0.01 1.5 -12.8 1.21 1.9 0.0 0.45 1.5 -10.0 2.45 2.3 1.7 0.01 1.0 -8.5 0.01 2.0 3.2 0.29 1.0 -7.0 1.19 1.1 5.3 0.47 2.0 -4.6 1.01 1.0 -1.9 0.78 2.1 2.9 0.35 2.5
59 Table 3.3: —Continued
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) HD 17505 -60.9 0.02 2.5 BD+60◦ 0586 -30.0 0.01 2.5 -50.9 0.20 2.5 -27.0 0.01 2.0 -45.3 0.52 2.5 -24.0 0.34 1.0 -38.6 0.27 2.3 -21.0 0.65 2.1 -33.4 0.49 2.3 -18.0 0.60 2.4 -27.3 0.46 2.3 -15.0 0.31 2.5 -23.2 0.64 1.1 -12.1 0.13 2.5 -19.9 0.50 1.1 -6.1 1.36 2.1 -16.7 1.02 1.8 -3.7 1.69 2.0 -13.1 0.49 1.5 -2.0 0.01 2.5 -10.5 0.21 1.0 0.1 1.87 2.1 -7.1 1.36 1.5 3.2 0.37 2.1 -4.6 1.12 1.0 5.2 0.01 2.0 -2.5 1.12 1.1 7.7 0.01 1.1 0.7 2.54 2.0 9.5 0.01 2.5 5.1 0.25 1.4 15.8 0.01 1.5
HD 17520 -48.7 0.39 2.5 HD 19243 -18.1 0.19 2.5 -44.8 0.66 2.5 -12.6 1.18 2.4 -37.5 0.17 1.5 -9.6 0.86 2.2 -32.7 0.48 2.5 -7.8 0.29 1.1 -28.1 0.28 2.0 -6.1 0.67 1.2 -23.0 0.61 2.3 -2.8 0.50 1.4 -19.0 0.57 1.9 -0.2 0.76 1.7 -14.6 1.01 2.5 2.2 0.28 1.3 -10.7 0.34 1.0 4.2 0.62 2.5 -7.7 2.48 1.1 10.5 0.24 2.0 -3.5 17.0 1.0 0.3 1.41 1.5 HD 19820 -19.7 1.10 2.3 3.4 0.67 2.5 -14.3 1.72 2.0 -8.8 0.98 2.5 -5.7 1.12 1.6 -2.6 0.82 1.2 0.0 1.01 1.5 2.9 0.42 1.3
60 Table 3.3: —Continued
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) HD 25443 -14.6 0.38 2.3 HD 192303 -15.7 0.07 2.5 -10.5 0.96 1.6 -9.2 0.61 2.5 -8.2 0.67 2.1 -5.8 0.01 1.0 -4.1 0.34 2.5 -3.9 0.26 1.9 1.1 1.56 1.4 -0.5 0.52 1.3 3.8 0.26 1.0 1.7 0.01 2.3 5.9 0.77 1.1 3.7 1.54 1.8 8.5 0.14 2.5 8.9 3.00 2.0 10.7 1.49 2.2 HD 25638 -15.9 0.38 1.8 12.3 0.10 2.2 -10.0 1.89 2.2 14.0 0.16 1.0 -8.2 0.17 1.6 16.0 0.28 2.3 -5.8 0.31 2.5 19.0 0.07 2.0 -3.1 0.13 1.3 22.0 0.12 1.5 -0.8 0.32 1.3 25.0 0.10 1.0 1.9 1.46 2.0 28.0 0.03 1.5 4.0 0.01 2.5 33.0 0.10 2.2 6.1 1.43 2.2 36.0 0.05 1.3 11.3 0.01 2.5 39.0 0.10 1.5
HD 45314 -17.0 0.16 2.5 HD 192641 -11.5 0.05 2.5 -11.0 0.37 1.9 -5.8 0.19 2.5 -5.3 0.35 2.5 -3.1 0.20 2.1 -2.6 0.90 1.8 -0.3 0.39 1.3 0.8 1.97 2.5 2.5 0.01 2.3 5.6 0.66 1.7 4.0 1.45 2.5 8.6 0.64 2.3 9.4 0.86 2.5 19.3 0.07 2.3 11.3 0.25 1.5 15.0 0.69 2.0 HD 46223 -18.0 0.11 2.5 18.0 0.16 1.0 -13.0 0.35 2.5 22.7 0.33 1.8 -11.0 0.13 2.3 -6.0 0.48 2.4 -3.0 0.10 1.8 1.0 0.75 2.0 5.0 1.59 2.3 8.1 1.52 1.5 11.0 1.07 1.0 13.5 1.74 2.5
61 Table 3.3: —Continued
Star Vlsr N b Star Vlsr N b (km s−1) (1012 cm−2) (km s−1) (km s−1) (1012 cm−2) (km s−1) HD 193793 -63.5 0.01 2.0 HD 197770 -12.5 0.09 2.5 -57.5 0.05 2.5 -5.8 0.64 1.9 -51.2 0.08 2.5 -3.1 0.97 1.0 -46.8 0.06 2.5 -1.1 1.16 1.5 -42.7 0.11 2.2 0.8 0.37 1.5 -39.1 0.13 1.9 2.5 1.32 1.7 -35.1 0.09 2.5 4.9 0.57 2.4 -30.1 0.18 2.3 10.3 0.41 2.0 -25.7 0.15 2.5 -21.8 0.24 2.5 HD 214419 -39.0 0.32 1.8 -19.3 0.25 1.3 -34.5 1.36 2.5 -16.4 0.38 1.7 -26.1 0.58 2.4 -13.0 0.27 1.8 -20.0 0.53 2.5 -10.4 0.23 1.5 -16.0 0.42 2.3 -6.3 0.55 2.4 -11.5 0.59 1.6 -4.1 0.38 1.8 -8.5 0.29 1.0 -1.8 1.29 1.4 -6.0 1.38 1.3 0.6 0.67 2.3 -4.0 1.72 1.5 5.1 1.37 2.2 -0.5 1.30 1.7 7.3 1.09 1.0 2.8 0.71 1.5 10.0 2.22 2.0 6.0 0.71 1.8 12.2 0.44 1.1 10.0 0.56 1.7 14.0 0.32 1.3 13.0 0.19 2.3 17.3 0.84 1.9 20.1 0.15 1.4 HD 217086 -35.0 0.03 1.3 22.6 0.19 2.2 -20.0 0.07 1.0 26.0 0.11 1.3 -15.5 0.13 2.5 27.8 0.12 1.6 -12.3 0.93 2.1 29.9 0.11 2.1 -6.3 1.19 2.1 34.5 0.04 2.1 -4.5 0.59 2.5 39.1 0.02 2.5 0.0 1.78 2.4 4.0 0.36 1.2 6.5 0.53 1.8 8.3 0.63 2.5
62 Table 3.4: Equivalent Width Measurements of CH, CH+, and CN from Mc- Donald Spectra
Star CH CH+ CN Vlsr Wλ Vlsr Wλ Vlsr Wλ (km s−1) (mA)˚ (km s−1) (mA)˚ (km s−1) (mA)˚ HD 108 -15.4 2.3(0.5) -16.4 1.7(0.6) ...... -11.7 1.6(0.4) ...... -9.6 1.7(0.2) ...... -4.8 3.3(0.5) ...... -5.1 4.3(0.9) ...... -0.8 3.7(0.9) ...... 0.8 1.4(0.6) 1.9 1.7(0.8) ...... 5.2 1.1(0.6) ...... 6.9 3.1(0.7) Total 10.6(2.8) 8.4(2.8) 7.2(1.6)
HD 5689 -20.3 1.7(0.9) -20.2 2.4(1.3) ...... -12.6 6.8(1.6) -13.2 8.6(2.1) -12.6 ≤10.3 -8.8 2.8(0.9) -8.0 9.6(2.5) ...... -1.4 3.3(2.1) -0.4 4.3(1.5) ...... Total 16.3(6.0) 25.8(7.7)
HD 12882 -5.4 3.0(0.5) ...... -2.4 3.4(0.6) -1.9 1.9(0.8) ...... 0.2 4.9(1.0) ...... 0.2 ≤6.5 2.1 4.3(0.7) 2.4 1.9(1.0) ...... 5.0 2.6(1.1) ...... Total 16.3(2.8) 6.0(2.9)
HD 13841 -41.6 3.3(0.2) -42.6 2.1(0.3) -41.6 ≤1.7 ...... -16.0 2.2(0.7) ...... -13.5 0.6(0.3) ...... -7.9 3.2(0.5) ...... -4.8 1.2(0.3) ...... -2.8 1.5(0.5) ...... -0.7 2.6(0.6) ...... Total 5.4(0.9) 11.8(2.9)
HD 13969 -15.0 1.4(0.5) -15.4 4.5(1.2) ...... -10.8 3.5(0.7) ...... -10.8 ≤4.6 ...... -8.0 4.4(1.4) ...... -1.7 4.4(1.5) ...... Total 5.5(1.2) 13.5(4.1)
63 Table 3.4: —Continued
Star CH CH+ CN Vlsr Wλ Vlsr Wλ Vlsr Wλ (km s−1) (mA)˚ (km s−1) (mA)˚ (km s−1) (mA)˚ HD 14053 ...... -13.8 4.1(0.9) ...... -8.6 7.2(1.1) -9.2 6.3(1.7) -8.6 ≤5.8 ...... -6.1 5.8(0.9) ...... -3.7 1.7(0.5) -3.5 4.0(1.1) ...... 2.9(0.5) -0.1 5.9(1.2) ...... Total 13.1(2.2) 27.5(5.9)
BD+ 56◦ 0501 -16.5 2.6(2.0) ...... -9.2 3.3(0.8) ...... -9.2 ≤6.7 -4.9 1.3(0.8) -4.7 6.3(2.4) ...... 0.9 4.4(1.3) ...... 16.0 3.2(1.2) ...... Total 8.0(4.0) 14.1(4.9)
BD+ 56◦ 0563 -12.3 1.5(0.8) ...... -8.4 2.7(1.7) -8.0 6.1(1.7) -8.4 ≤4.4 -5.5 3.0(1.8) -5.5 2.4(2.0) ...... -3.5 3.7(1.4) ...... -1.3 1.3(0.8) ...... Total 8.8(5.1) 12.8(5.7)
HD 14443 ...... -10.1 4.4(0.9) ...... -6.2 2.7(1.1) -5.3 5.1(1.1) -6.2 ≤3.1 -1.0 1.4(0.5) 0.0 6.0(1.0) ...... Total 4.4(1.6) 15.5(3.0)
HD 14476 -7.9 2.0(1.0) ...... -4.3 2.4(1.0) -6.0 5.6(3.3) ...... -0.6 1.3(1.1) -0.4 5.3(3.3) -0.6 ≤7.5 Total 5.5(3.2) 11.6(6.6)
BD+ 56◦ 0578 -23.2 1.6(0.9) ...... -12.4 3.6(1.3) ...... -12.4 ≤8.1 -6.6 1.7(0.9) -7.1 3.7(4.5) ...... -4.2 2.5(1.7) ...... Total 7.2(3.1) 6.9(6.4)
64 Table 3.4: —Continued
Star CH CH+ CN Vlsr Wλ Vlsr Wλ Vlsr Wλ (km s−1) (mA)˚ (km s−1) (mA)˚ (km s−1) (mA)˚ HD 14947 -11.0 7.3(0.9) -11.5 7.6(1.2) ...... -7.7 7.6(0.6) -7.3 4.2(1.0) -7.7 ≤4.7 -4.8 4.7(0.9) -3.9 4.5(0.9) ...... Total 19.7(2.6) 16.9(3.1)
HD 15629 -5.1 3.8(1.1) -5.9 3.4(1.7) -5.1 ≤9.0 ...... -2.5 3.7(2.1) ...... -0.3 3.0(0.9) ...... 6.7 3.4(1.8) ...... Total 8.3(2.0) 10.6(5.6)
HD 16691 -46.9 6.0(1.7) ...... -12.7 7.3(2.4) -12.1 6.3(2.4) -12.7 ≤16.3 ...... -4.1 17.6(5.6) ...... 4.4 2.8(1.6) 4.9 10.7(2.9) ...... Total 15.4(5.9) 37.9(11.0) HD 17505 -6.8 4.6(0.5) -5.4 17.6(0.9) ...... -4.3 5.9(0.4) -3.4 6.2(1.4) -4.3 ≤3.8 -1.9 3.6(0.3) ...... 1.1 2.0(0.3) 1.3 2.4(1.0) ...... 5.4 1.1(0.3) ...... Total 18.1(2.0) 27.1(4.5)
HD 17520 -5.9 2.8(0.6) ...... -2.5 3.7(0.7) -1.8 5.0(1.0) -2.5 ≤5.5 ...... 1.6 3.0(1.0) ...... Total 7.3(1.3) 9.2(2.1)
BD+ 60◦ 0586 -6.4 1.8(0.4) -6.1 4.6(1.9) ...... -3.6 1.4(0.7) ...... -2.2 4.2(0.8) -2.2 0.8(0.8) -2.2 ≤4.5 ... 4.7(1.0) 0.3 2.8(0.8) ...... 3.1 4.3(1.0) ...... Total 11.5(2.2) 14.0(5.8)
65 Table 3.4: —Continued
Star CH CH+ CN Vlsr Wλ Vlsr Wλ Vlsr Wλ (km s−1) (mA)˚ (km s−1) (mA)˚ (km s−1) (mA)˚ HD 19243 ...... -11.9 2.4(0.8) ...... -10.0 2.4(0.6) -9.5 3.4(0.5) ...... -7.7 3.8(0.6) -6.3 1.7(0.5) ...... 0.5 3.9(0.6) ... 1.4(1.1) 0.5 ≤2.2 ...... 2.6 6.8(1.2) ...... Total 8.0(1.8) 15.5(5.0)
HD 19820 ...... -19.1 4.1(1.8) ...... -14.9 12.0(0.8) -15.0 4.7(1.2) -14.9 ≤6.6 ...... -5.8 2.4(0.7) ...... -3.4 2.3(0.7) -2.8 3.3(0.7) ...... -0.8 3.0(0.4) ...... 0.8 3.6(1.0) ...... 2.6 2.2(0.6) ...... Total 21.3(3.1) 17.8(5.6)
HD 25443 ...... -14.2 7.6(1.1) ...... -10.9 3.1(0.7) -10.9 4.7(0.9) ...... -8.1 3.6(0.5) -7.5 5.4(0.9) -8.1 ≤4.0 1.8 2.5(0.5) ...... 6.4 2.5(0.4) ...... Total 13.0(2.1) 18.3(2.9)
HD 25638 -16.1 2.8(0.4) -16.1 2.3(1.2) ...... -10.2 10.1(0.7) -10.0 8.6(1.1) ...... 2.2 7.6(0.6) ...... 2.4 8.6(1.0) 6.5 7.6(0.4) 5.3 4.2(1.0) 5.7 5.1(0.9) Total 30.5(2.2) 14.8(3.9) 13.9(1.9)
HD 45314 ...... -5.1 1.4(0.7) ...... 1.7 4.4(1.1) ...... 1.7 ≤4.9 5.1 4.2(0.7) ...... 8.3 1.7(1.1) ...... Total 8.6(1.8) 2.9(1.8)
66 Table 3.4: —Continued
Star CH CH+ CN Vlsr Wλ Vlsr Wλ Vlsr Wλ (km s−1) (mA)˚ (km s−1) (mA)˚ (km s−1) (mA)˚ HD 46223 1.4 1.7(0.4) 0.2 2.5(1.3) ...... 8.2 5.5(0.8) 8.1 5.7(1.8) ...... 11.3 7.4(0.6) 10.8 12.8(1.2) 11.3 ≤3.6 Total 15.8(1.9) 21.3(5.3)
HD 192303 3.4 1.8(1.0) 4.1 11.0(2.0) ...... 8.9 4.4(1.8) ...... 8.9 ≤8.6 Total 7.4(2.8) 12.0(2.0)
HD 192641 ... 2.1(0.8) ...... 4.0 7.0(0.9) 3.9 16.6(1.6) 4.0 ≤4.0 ...... 9.4 5.0(1.4) ...... Total 9.8(2.0) 22.3(3.4)
HD 193793 -3.5 2.1(0.5) -3.2 4.6(1.0) ...... 6.2 5.0(0.5) 5.6 14.9(0.9) 6.2 ≤1.8 10.1 4.9(0.6) 10.1 22.6(1.0) ...... 13.9 5.0(0.5) ...... 18.2 0.8(0.2) 17.2 4.0(0.6) ...... Total 18.2(2.5) 47.1(4.3)
HD 197770 ...... -3.5 0.8(0.3) ...... -1.1 20.8(0.3) -0.7 1.5(0.4) -1.4 15.6(0.6) 2.5 4.6(0.7) ...... Total 26.0(1.7) 2.1(0.7) 16.8(0.6) HD 214419 -18.8 5.4(0.4) ...... -19.5 6.7(1.6) ...... -7.9 2.9(0.8) ...... -5.6 3.7(0.8) -5.6 -3.0(1.8) ...... -3.9 11.2(1.0) ...... -0.4 2.3(0.7) -0.9 0.9(0.8) ...... 6.9 7.3(1.2) 6.3 1.8(1.0) ...... Total 20.0(3.4) 13.8(2.6) 7.4(1.6)
67 Table 3.4: —Continued
Star CH CH+ CN Vlsr Wλ Vlsr Wλ Vlsr Wλ (km s−1) (mA)˚ (km s−1) (mA)˚ (km s−1) (mA)˚ HD 217086 -21.6 2.5(0.5) -20.2 3.4(1.1) ...... -15.7 11.1(0.8) -14.1 3.8(0.7) -15.7 ≤5.8 -11.7 7.3(0.8) -11.4 6.2(1.0) ...... -4.5 3.2(1.1) -6.7 5.9(0.8) ...... 0.1 8.4(0.9) ...... 6.5 1.6(0.5) 6.4 4.7(1.0) ...... Total 26.4(4.5) 32.1(5.9)
68 Table 3.5: UVES Measurements for Southern Hemisphere Sight Lines
Star CH CH+ CN a Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 92964 ...... -11.9 0.5(0.1) 2.9 ...... -8.8 2.2(0.1) 2.8 ...... -9.8 0.03(0.01)/0.015b 0.5 -5.8 6.3(0.1) 1.5 -5.8 9.0(0.1) 2.8 -5.8 0.17(0.01)/0.09(0.01) 1.1 -1.0 4.1(0.1) 2.2 0.3 0.4(0.1) 3.3 0.2 0.27(0.01)/0.12(0.01) 1.0 Total 12.6 9.9 0.695
HD 97253 ...... -18.3 7.7(0.4) 3.0 ...... -15.2 13.4(0.3) 2.0 -15.4 2.1(0.3) 1.4 -14.1 2.26(0.06)/1.03(0.08) 0.9 -7.0 4.2(0.3) 0.8 -5.9 4.6(0.4) 2.5 -7.3 0.11(0.05)/0.055b 0.8 -1.4 1.7(0.3) 0.9 -0.4 5.1(0.3) 1.5 ......
69 Total 19.3 19.5 3.455
HD 108927 2.4 14.1(0.2) 1.5 2.6 5.2(0.3) 3.3 1.5 1.42(0.03)/0.61(0.04) 1.0 Total 14.1 5.2 2.030
HD 149404 -10.8 5.9(0.2) 1.2 -10.1 5.9(0.2) 2.5 -11.3 0.48(0.06)/0.20(0.04) 0.3 0.4 8.4(0.2) 3.0 0.3 21.1(0.3) 3.5 ...... 7.7 14.2(0.1) 0.4 7.3 21.5(0.2) 1.2 7.9 2.57(0.73)/1.04(0.10) 0.3 Total 28.5 48.5 4.290 Table 3.5: —Continued
Star CH CH+ CN a Vlsr N b Vlsr N b Vlsr N b (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) (km s−1) (1012cm−2) (km s−1) HD 151932 0.2 4.7(0.1) 3.0 0.4 11.3(0.1) 2.3 0.1 0.05(0.02)/0.025b 1.1 6.9 26.0(0.1) 2.2 6.1 3.9(0.1) 2.9 7.1 1.77(0.02)/0.85(0.03) 2.1 Total 30.7 15.2 2.695
HD 168076 ...... -2.5 4.5(0.2) 2.1 ...... 0.6 9.1(0.2) 3.6 ...... 1.4 1.23(0.06)/0.52(0.05) 0.5 3.4 9.7(0.2) 3.2 3.6 6.1(0.2) 2.2 ...... 70 7.5 2.4(0.1) 1.2 7.1 4.2(0.2) 2.1 6.5 0.22(0.03)/0.110b 1.0 14.6 1.0(0.2) 3.5 ...... 19.8 2.9(0.1) 1.9 20.0 4.9(0.2) 2.8 ...... 22.9 3.8(0.2) 3.0 22.5 5.2(0.2) 2.3 ...... 27.6 0.6(0.1) 1.2 27.6 1.0(0.2) 1.7 ...... Total 29.5 25.9 2.080
HD 170740 ...... -1.1 0.5(0.1) 2.8 ...... 2.9 4.4(0.1) 1.0 ...... 6.5 16.0(0.1) 1.0 5.1 16.5(0.1) 3.1 7.1 5.42(0.45)/2.41(0.11)d 0.5 Total 20.4 17.0 7.830 aColumn densities for CN are given as: N (J = 0)/N (J = 1). bThe J = 1 lines of CN were not detected, N (J = 1) is assumed to be 50% of N (J = 0). cHD 170740 has a CN J = 2 column density of 0.11(0.02). Table 3.6: McDonald Results for the GOT C+ Project Sight Lines
Star Ca II K I Ca I CH+ CH CN
11 11 11 11 11 a 11 Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) HD 165783 -21.5 0.6 2.5 ...... -15.9 2.4 2.0 ...... -15.3 58.0 2.4 ...... -13.0 2.0 2.0 ...... -9.8 12.2 2.2 ...... -10.5 0.10 1.6 -9.1 51.9 2.5 ...... -6.4 4.2 1.4 ...... -3.6 4.3 2.5 ...... -4.3 26.0 1.9 ...... 1.1 6.4 2.5 ...... 4.9 40.3 1.6 ...... 6.0 0.10 1.6 5.0 228.0 2.5 4.1 67.3 1.5 ...... 7.7 8.4 1.6 ...... 6.9 171.0 1.4 7.3 138.0 1.1 7.3 ≤ 112 1.0 9.9 7.6 2.2 ...... 13.5 2.1 1.6 ...... 16.6 2.2 2.0 ...... 20.4 4.0 2.2 ...... 23.8 2.9 2.5 ...... 28.5 4.7 2.5 ...... 34.8 2.2 1.6 ...... 39.0 4.4 2.4 ...... 45.2 0.5 1.4 ......
71 HD 165918 -21.2 0.7 2.5 ...... -16.0 2.1 2.3 ...... -11.2 3.2 2.3 ...... -8.8 1.9 2.2 ...... -7.7 43.1 3.3 ...... -3.5 1.8 2.5 ...... 0.0 5.1 2.5 ...... 4.2 17.6 2.3 ...... 5.0 ≤ .057 1.6 5.1 128.5 2.3 4.9 70.2 2.5 5.1 ≤ 79.7 1.0 10.7 1.9 2.5 ...... 16.0 1.9 2.5 ...... Table 3.6: —Continued
Star Ca II K I Ca I CH+ CH CN
11 11 11 11 11 a 11 Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) HD 168607 -45.5 1.0 2.5 ...... -39.7 2.2 2.5 ...... -33.6 3.1 2.5 ...... -28.4 4.5 2.5 ...... -23.0 4.1 2.5 ...... -18.2 6.0 2.1 ...... -13.4 5.9 1.8 ...... -9.8 5.8 1.8 ...... -5.3 10.0 2.4 ...... -0.8 7.4 1.7 ...... 3.4 118.8 1.4 ...... 1.8 31.5 1.0 ...... 6.7 11.5 1.8 ...... 5.5 ≤.076 1.6 5.5 177.6 2.9 7.0 85.1 2.5 ...... 72 9.9 11.8 2.4 ...... 15.0 16.4 2.2 ...... 15.2 333.0 2.1 14.5 199.3 2.1 ...... 18.7 9.0 1.4 ...... 18.3 47.7 1.0 18.5 26 1.7 22.1 25.3 2.0 ...... 20.8 360.0 3.4 22.9 344.9 1.6 ...... 26.0 10.9 1.8 ...... 27.2 154.0 3.5 ...... 30.7 4.4 1.6 ...... 35.4 3.4 2.0 ...... 40.5 3.3 2.5 ...... 46.2 2.4 2.5 ...... 52.6 1.4 2.0 ...... 59.2 1.9 2.1 ...... 65.6 0.7 2.5 ......
HD 167498 -14.7 1.1 2.5 ...... -10.8 2.4 1.6 ...... -6.0 1.4 1.8 ...... -6.5 21.2 1.6 ...... 0.6 12.0 1.8 ...... 1.6 0.10 1.6 0.8 99.4 2.7 0.1 43.4 1.7 0.1 ≤32 1.0 4.1 6.2 2.3 ...... 4.4 67.9 3.2 5.1 39.0 2.5 ......
HD 167812 1.0 12.6 1.8 ...... 2.2 0.10 1.6 1.6 101.8 2.0 1.2 121.5 1.8 1.2 ≤ 65 1.0 4.1 4.9 1.8 ...... 5.4 16.7 0.9 ...... Table 3.6: —Continued
Star Ca II K I Ca I CH+ CH CN
11 11 11 11 11 a 11 Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) HD 169754 -11.6 1.6 1.5 ...... -7.3 2.1 2.4 ...... -3.1 6.9 1.9 ...... -3.1 64.7 3.5 -1.0 64.0 1.2 -2.0 66 1.7 1.6 59.6 2.0 ...... 2.5 0.20 1.6 1.9 87.5 1.1 2.5 171.3 0.8 1.5 170 1.2 5.5 500.0 1.8 ...... 4.6 74.1 1.1 5.8 281.0 1.3 5.5 340 0.7 12.2 22.7 2.3 ...... 16.3 0.20 1.6 15.7 46.1 2.7 ...... 19.0 500.0 2.5 ...... 22.9 0.10 1.2 21.4 58.7 2.8 ...... 26.4 8.7 2.5 ...... 32.3 4.8 2.5 ...... 38.2 8.7 2.5 ...... 42.6 22.5 2.5 ...... 48.3 7.1 2.5 ...... 60.3 4.5 2.5 ...... 63.5 8.2 1.4 ......
HD 174509 -9.1 1.3 2.5 ...... -5.1 2.8 1.0 ...... 73 0.1 8.9 2.5 1.0 2.5 1.5 1.0 0.09 0.4 0.3 38.4 1.2 0.7 12.4 1.1 ...... 2.8 6.6 1.7 3.8 1.8 2.3 3.2 0.07 1.9 ...... 3.5 20.7 1.7 3.2 5.5 0.3 7.8 23.2 2.2 8.4 9.6 0.7 8.9 0.16 1.2 7.6 20.7 2.5 8.4 132.5 2.5 7.4 11.0 0.3 12.2 124.1 1.3 11.5 2.8 2.1 11.6 0.15 2.5 10.5 39.4 1.4 12.0 41.4 1.0 11.3 14.0 1.1 15.7 15.1 2.5 ...... 14.9 0.18 1.5 ......
BD+49◦ 3482 -4.4 0.4 1.3 ...... 1.6 1.5 1.4 1.2 1.2 1.1 ...... 1.5 21.3 2.0 ...... 3.8 2.0 2.0 ...... 3.7 22.2 2.5 ...... 6.8 1.5 2.5 ......
BD+49◦ 3484 3.5 3.7 2.3 ...... 2.8 57.5 3.5 ...... 8.5 3.8 2.0 ...... 10.3 1.6 2.0 ...... 10.3 75.1 1.9 10.6 ≤ 1.9 1.0 13.0 1.2 2.5 ......
HD 240179 -10.3 1.0 2.5 ...... -9.1 0.09 1.5 -9.9 69.6 2.1 ...... -6.9 9.6 2.4 -6.6 1.0 2.2 -6.0 0.13 2.0 -7.2 43.4 2.5 -6.6 30.9 1.6 ...... -3.5 44.2 2.3 -3.7 2.1 2.0 -3.0 0.19 1.9 -3.3 73.2 2.3 ...... -1.2 17.2 2.2 -1.5 5.4 1.4 ...... -1.1 24.7 3.3 -1.5 40.2 2.0 ...... 0.2 1.9 2.0 ...... 0.3 66.1 2.0 0.3 ≤ 8.4 1.0 4.0 6.5 2.5 3.5 0.5 2.1 3.9 0.05 1.0 ...... 7.6 3.5 2.5 ...... 11.9 1.6 2.5 ...... 17.5 1.0 2.4 ...... Table 3.6: —Continued
Star Ca II K I Ca I CH+ CH CN
11 11 11 11 11 a 11 Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) HD 240183 -12.5 0.6 2.5 ...... -12.4 0.09 1.5 ...... -9.9 2.3 1.5 ...... -5.8 9.6 1.9 -5.3 0.4 1.0 ...... -5.3 49.5 3.5 ...... -1.9 20.8 2.5 -1.7 4.4 1.3 -3.3 0.12 1.5 -1.6 42.7 2.2 -1.3 61.1 2.5 -1.3 ≤ 5.5 1.0 2.1 5.9 2.5 ...... 6.9 4.3 2.5 ...... 12.1 1.8 2.5 ...... 18.0 2.5 1.8 ......
HD 35652b -32.8 1.0 2.5 ...... -26.9 1.7 2.4 ...... -22.3 1.3 2.2 ......
74 -18.4 2.2 1.2 ...... -14.4 4.4 1.5 -14.5 0.3 2.5 ...... -15.6 23.9 2.9 ...... -11.3 6.1 1.3 -11.1 0.4 1.8 ...... -9.0 6.0 1.2 ...... -9.0 19.2 1.5 ...... -6.1 9.8 2.5 -6.2 0.1 1.5 ...... -6.1 35.1 3.5 ...... -2.5 6.5 2.5 -2.0 0.3 1.0 ...... 0.0 3.3 2.5 ...... 3.7 20.9 1.7 3.8 9.2 1.4 3.1 0.14 1.8 3.7 147.8 2.9 3.7 117.2 2.3 3.7 ≤ 6.4 1.0 6.2 0.6 2.5 ...... 6.1 135.8 2.2 ...... 8.7 4.4 1.7 ...... 13.2 0.9 2.5 ......
HD 47073 3.8 2.2 2.5 ......
HD 260737 -15.6 1.6 2.5 ...... -8.6 0.9 2.0 ...... -2.9 1.9 2.3 ...... 1.8 2.7 2.2 ...... 5.3 1.7 1.8 ...... 7.5 0.8 2.3 7.2 0.3 2.2 ...... 11.1 0.9 2.3 ...... Table 3.6: —Continued
Star Ca II K I Ca I CH+ CH CN
11 11 11 11 11 a 11 Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b Vlsr N /10 b (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) (km s−1) (cm−2) (km s−1) HD 55469 1.6 3.8 1.4 1.8 1.2 1.4 ...... 3.8 3.5 2.3 ...... 10.3 1.6 2.3 ......
75 HD 55981 -7.3 0.5 2.1 ...... -2.3 2.5 2.3 ...... 2.8 2.7 2.5 ...... 5.7 3.2 2.3 ...... 9.2 4.6 2.5 ...... 14.8 4.5 2.5 14.9 1.3 1.0 ...... 14.5 125.5 2.2 ...... 20.9 2.3 2.5 ......
HD 60146b ...... 0.0 0.4 1.0 ......
aThe column density for CN includes J = 0 and J = 1 for detections. For upper limits, the column density for CN includes J = 0 (at the maximum depth at the largest CH component with b = 1 km s−1) and includes the predicted J = 1 column density, 50% of the J = 0 column density. bNo additional data. 3.2.1 Comparison with Literature Values
Table 3.7 is a comparison of column density and equivalent width for the sight
lines from the current data set with available results from recent work. The entries
are grouped by sight line and subgrouped by the species and line, given as the first
three columns. The next two columns are the measurements of the current data set.
The total equivalent width of the observed line is denoted Wλ. The next two columns are the total equivalent width and total column density from previous work. These
columns are followed by the instrument, location or facility, and the stated resolution
(when available) from each of the corresponding references listed in the last column.
Uncertainties in the values from previous work are also given when available.
The measured equivalent widths toward HD 15629 in CH and CH+ are consider- ably smaller then those quoted in Kre lowski et al. (1999). However, the CH+ portion of our spectrum is quite noisy and given the lack of error bars from Kre lowski et al.
(1999), it is difficult to make an accurate comparison.
The current sample has several sight lines in common with the Ca II measurements from Megier et al. (2009). The measured equivalent width toward HD 45314 is smaller than that found in Megier et al. (2009) but agrees at the 2 sigma level. However, the measured column densities of Ca II are consistent with the values found in Megier et al. (2009). The measured Ca II equivalent widths and column densities toward
HD 46223 and HD 193793 agree with the values found in Megier et al. (2009). The measured equivalent widths toward HD 217086 in Ca II are only similar to those of
Megier et al. (2005) at the 3 sigma level, and the measured equivalent widths in CH also only agree with the values of Megier et al. (2005) at the 2 sigma level. Values from Federman et al. (1994) of equivalent widths toward HD 217086 were derived with a curve of growth treatment. The values from Federman et al. (1994) for CH and CN in this same sight line are similar to our measured values. Although there is
76 no detection of CN in our measurements, our upper limit is also consistent with the value from Federman et al. (1994).
There are two common lines of sight between the current sample and those from
Weselak et al. (2008) and S lyk et al. (2008). It is worth noting that measurements of Weselak et al. (2008) and S lyk et al. (2008) come from the same data set. Our measured equivalent widths toward HD 45314 agree with those from both Weselak et al. (2008) and S lyket al. (2008). The column densities are within one sigma of the values from Weselak et al. (2008), and within 2 sigma of the S lyket al. (2008) values.
Although there is no detection of CN toward HD 45314 in our measurements, our upper limits also agree with the quoted values from both Weselak et al. (2008) and
S lyket al. (2008). The measured CH equivalent widths and column densities toward
HD 46223 are consistent with the values of Weselak et al. (2008). However, S lyket al.
(2008) quote an equivalent width significantly different from our values and those of
Weselak et al. (2008). Although there is no detection of CN in our measurements toward HD 46223, our upper limits are also consistent with the values from both
Weselak et al. (2008) and S lyket al. (2008).
The CN values of equivalent width and column density from Weselak et al. (2008) toward HD 25638 are inconsistent with our values, but the equivalent width of CN from Kos & Zwitter (2013) is similar to our value. Some of the inferred CN col- umn density values toward HD 170740 from Weselak et al. (2009b) disagree because there were significant optical depth effects present in the main component, which was excluded in their analysis. Kre lowski et al. (2012) notes severe saturation of the observed CN line, values with and without the saturation corrections are provided in the table. In general, there is good agreement between the current data set and the values from previous work in the literature, with few exceptions.
77 Table 3.7: Total Column Density and Wλ Comparison
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD 15629 CH+ 4232.548 10.7(5.7) 13.5 23 ... OHP ELODIE 42000 Kre lowski et al. 1999 CH 4300.313 8.3(1.7) 10.3 17 ... OHP ELODIE 42000 Kre lowski et al. 1999
HD 25638 Ca II 3933.663 245(5) ... 230.6(2.8) ... Padua echelle 23000 Kos et al. 2013 CH+ 4232.548 14.8(3.3) 21.6 13.6(6.7) ... Padua echelle 23000 Kos et al. 2013 4232.548 ≤ 20 13.31(3.47) k BOAO ... m ∼ 45000 n Weselak et al. 2014 B CH 4300.313 31.0(2.1) 21.6 26.0(3.0) 31.4(3.62) Terskol MAESTRO 80000 Weselak et al. 2008 4300.313 ≤ 20 31.40(7.47) j BOAO ... m ∼ 45000 n Weselak et al. 2014 B 4300.313 27.0(1.0) ... Padua echelle 23000 Kos et al. 2013 3886.410 9.5(0.3) ... Terskol MAESTRO 80000 Weselak et al. 2008 78 CN 3874.602 13.7(1.8) 5.8 1.0(0.2) 0.33 e Terskol MAESTRO 80000 Weselak et al. 2008 3874.602 1.0(0.2) 0.22(0.04) i Terskol MAESTRO 80000 Weselak et al. 2008 3874.602 17.8(2.2) ... Padua echelle 23000 Kos et al. 2013
HD 45314 Ca II 3933.663 227(13) 5.5 277(14) 6.23 ESO m FEROS m 48000 m Megier et al. 2009 CH+ 4232.548 2.9(1.8) 3.8 ≤ 20 7.18(1.97) k ESO FEROS 48000 Weselak et al. 2014 B CH 4300.313 9.7(1.7) 11.2 13.7(2.0) 16.54(2.41) ESO FEROS 48000 Weselak et al. 2008 4300.313 ≤ 20 15.44(0.60) j ESO FEROS 48000 Weselak et al. 2014 B 4300.313 11.74(0.4) 18.00 ESO FEROS 48000 S lyket al. 2008 3886.410 4.1(0.1) ... ESO FEROS 48000 Weselak et al. 2008 3886.410 3.85(0.4) ... ESO FEROS 48000 S lyket al. 2008 CN 3874.602 ≤ 7.1 a ≤ 1.8 6.5(0.3) 2.19(0.13) ESO FEROS 48000 Weselak et al. 2008 3874.602 6.37(0.6) 2.34f ESO FEROS 48000 S lyket al. 2008 3873.994 2.3(0.2) ... ESO FEROS 48000 Weselak et al. 2008 3873.994 2.44(0.4) ... ESO FEROS 48000 S lyket al. 2008 3875.758 1.8(0.1) ... ESO FEROS 48000 Weselak et al. 2008 3875.758 2.26(0.5) ... ESO FEROS 48000 S lyket al. 2008 Table 3.7: —Continued
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD 46223 Ca II 3933.663 294(11) 8.2 297(15) 6.63 ESO m FEROS m 48000 m Megier et al. 2009 CH+ 4232.548 21.3(4.3) 20.6 ≤ 20 28.93(1.62) k ESO FEROS 48000 Weselak et al. 2014 B CH 4300.313 16.0(1.8) 20 18.6(1.0) 22.46(1.21) ESO FEROS 48000 Weselak et al. 2008 4300.313 ≤ 20 21.88(0.60) j ESO FEROS 48000 Weselak et al. 2014 B 4300.313 7.51(0.4) ... ESO FEROS 48000 S lyket al. 2008 3886.410 7.4(0.2) ... ESO FEROS 48000 Weselak et al. 2008 3890.217 5.2(0.2) ... ESO FEROS 48000 Weselak et al. 2008 CN 3874.602 ≤ 6.7a ≤ 1.3 4.7(0.4) 1.43(0.15) ESO FEROS 48000 Weselak et al. 2008 3874.602 4.63(0.8) 1.55f ESO FEROS 48000 S lyket al. 2008 3873.994 1.2(0.2) ... ESO FEROS 48000 Weselak et al. 2008 3873.994 1.33(0.3) ... ESO FEROS 48000 S lyket al. 2008 3875.758 0.8(0.1) ... ESO FEROS 48000 Weselak et al. 2008 79 3875.758 0.84(0.2) ... ESO FEROS 48000 S lyket al. 2008
HD 197770 Ca II 3933.663 193 ... 186(1) ... CFHT ESPaDOnS 64000 Cox et al. 2011 CH+ 4232.548 2.1(0.7) 2.7 3.0(0.3) ... CFHT ESPaDOnS 64000 Cox et al. 2011 CH 4300.313 26.0(1.4) 40.8 26.6(0.1) ... CFHT ESPaDOnS 64000 Cox et al. 2011 CN 3874.602 14.9 5.9 20.0(0.3) ... CFHT ESPaDOnS 64000 Cox et al. 2011 3873.998 4.1(0.2) ... CFHT ESPaDOnS 64000 Cox et al. 2011
HD 193793 Ca II 3933.663 504(47) 13 549(48.0) 11.6 ... m ... m ... m Megier et al. 2009
HD 217086 Ca II 3933.663 269(12) 6.2 303.2(1.5) ...... m ... m ... m Megier et al. 2005 3933.663 303(15) 7.03(1.13) ... m ... m ... m Megier et al. 2009 3933.663 262.2(4.6) ... Padua echelle 23000 Kos et al. 2013 CH+ 4232.548 32.1(5.5) 40.0 39.7(2.0) ... Padua echelle 23000 Kos et al. 2013 CH 4300.313 29.0(3.7) 36.0 41.5(6.2) ...... m ... m ... m Megier et al. 2005 4300.313 31.0(3.4) ... Padua echelle 23000 Kos et al. 2013 4300.313 26.8 c 55 Kitt Peak Cassegrain 35000 Federman et al. 1994 CN 3874.602 ≤ 8a ≤2.4 4.2 d 1.55 Kitt Peak Cassegrain 35000 Federman et al. 1994 3874.602 10.7(0.5) ... Padua echelle 23000 Kos et al. 2013 Table 3.7: —Continued
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD 92964 Ca II 3933.663 ...... 513(26) 6.5(0.74) ... m ... m ... m Megier et al. 2009 3933.663 512.7(3.8) ...... m ... m ... m Megier et al. 2005 CH+ 4232.548 8.3(0.1) 10.0 8.06(0.02) ... ESO VLT UVES 75000 Stahl et al. 2008 4232.548 8.22(0.09) 9.51(0.10) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 4232.548 ≤ 20 9.647(0.12) k ESO VLT UVES 80000 Weselak et al. 2014 B 3957.689 4.46(0.02) ... ESO VLT UVES 75000 Stahl et al. 2008 3957.689 4.50(0.09) 9.49(0.19) ESO VLT m UVES m 80000 m Weselak et al. 2014 A CH 4300.313 ... 12.6 10.4(0.2) ...... m ... m ... m Megier et al. 2005 4300.313 ≤ 20 12.69(0.16) j ESO VLT UVES 80000 Weselak et al. 2014 B 80 4300.313 10.11(0.05) 12.52(0.06) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 3886.410 2.92(0.08) 6.86(0.12) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 3890.217 1.84(0.08) 6.55(0.28) ESO VLT m UVES m 80000 m Weselak et al. 2014 A CN 3874.602 ... 0.71 1.84(0.07) 0.61438 e ESO VLT UVES 80000 Kre lowski et al. 2012 3874.602 ... 0.47 1.84(0.07) 0.4048 g ESO VLT UVES 80000 Kre lowski et al. 2012 3874.602 1.84(0.07) 0.4119(0.0127) h ESO VLT UVES 80000 Kre lowski et al. 2012 3873.994 0.60(0.10) 0.20248(0.0184) h ESO VLT UVES 80000 Kre lowski et al. 2012 T01 2.88(0.18) ... g ESO VLT UVES 80000 Kre lowski et al. 2012 T01 2.86(0.11) ... h ESO VLT UVES 80000 Kre lowski et al. 2012
HD 97253 CH+ 4232.548 16.2(0.6) 19.5 18.15(1.08) 21.01 ESO VLT UVES 80000 Weselak et al. 2009 A CH 4300.313 10.2(0.1) 19.1 11.20(0.5) 15.99 ESO VLT UVES 80000 S lyket al. 2008 3886.410 3.42(0.4) ... ESO VLT UVES 80000 S lyket al. 2008 CN 3874.602 9.0(0.3) 3.3 8.72(0.6) 3.70 ESO VLT UVES 80000 S lyket al. 2008 3873.994 4.15(0.3) ... ESO VLT UVES 80000 S lyket al. 2008 3875.758 1.35(0.6) ... ESO VLT UVES 80000 S lyket al. 2008 Table 3.7: —Continued
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD149404 Ca II 3933.663 ...... 340(17) 7.32(1.19) ESO m FEROS m 48000 m Megier et al. 2009 3933.663 339.8(2.3) ... ESO m FEROS m 48000 m Megier et al. 2005 CH+ 4232.548 37.0(0.3) 49.0 37.5 48 ESO CAT CES 60000 Gredel et al. 1993 4232.548 37.04(0.30) 42.87(0.35) ... m ... m ... m Weselak et al. 2014 A 4232.548 ≤ 20 43.63(2.20) k BOAO ... m ∼ 45000 n Weselak et al. 2014 B 4232.548 ≤ 20 49.50(0.46) k ESO VLT UVES 80000 Weselak et al. 2014 B 3957.689 21.9(3.0) 11(2) ESO CAT CES 60000 Gredel et al. 1993 3957.689 21.12(0.32) 44.55(0.68) ... m ... m ... m Weselak et al. 2014 A CH 4300.313 21.5(0.2) 28.2 23.3 31 ESO CAT CES 60000 Gredel et al. 1993 4300.313 22.81(0.29) 28.66(0.36) ... m ... m ... m Weselak et al. 2014 A 4300.313 ≤ 20 31.50(1.02) j BOAO ... m ∼ 45000 n Weselak et al. 2014 B 4300.313 ≤ 20 28.63(0.32) j ESO VLT UVES 80000 Weselak et al. 2014 B 4300.313 10.80(0.2) 14.82 ESO FEROS 48000 S lyket al. 2008 4300.313 24.5(0.5) ... ESO m FEROS m 48000 m Megier et al. 2005
81 4300.313 22.5(0.9) 27.17(1.09) ESO FEROS 48000 Weselak et al. 2008 3886.410 5.46(0.29) 12.91(0.67) ... m ... m ... m Weselak et al. 2014 A 3886.410 3.13(0.2) ... ESO FEROS 48000 S lyket al. 2008 3886.410 5.7(0.1) ... ESO FEROS 48000 Weselak et al. 2008 3890.217 3.95(0.33) 14.17(1.17) ... m ... m ... m Weselak et al. 2014 A 3890.217 3.2(0.1) ... ESO FEROS 48000 Weselak et al. 2008 Table 3.7: —Continued
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD 149404 CN 3874.602 9.2(0.2) 4.1 6.1(0.7) 1.9 e ESO CAT CES 60000 Gredel et al. 1991 3874.602 6.1(0.7) 1.5(0.2) ESO CAT CES 60000 Gredel et al. 1991 3874.602 7.73(0.43) ... ESO CAT CES 120000 Palazzi et al. 1992 3874.602 7.81(0.2) 2.89 ESO FEROS 48000 S lyket al. 2008 3874.602 8.7(0.8) 2.74(0.31) ESO FEROS 48000 Weselak et al. 2008 3874.602 9.85(0.41) 3.3473 e ESO VLT UVES 80000 Kre lowski et al.2012 3874.602 9.85(0.41) 2.1668(0.091) g ESO VLT UVES 80000 Kre lowski et al.2012 3874.602 9.85(0.41) 2.2971(0.044) h ESO VLT UVES 80000 Kre lowski et al.2012 3873.994 0.7(0.4) 0.4(0.3) ESO CAT CES 60000 Gredel et al. 1991 3873.994 2.48(0.60) ... ESO CAT CES 120000 Palazzi et al. 1992 3873.994 2.68(0.2) ... ESO FEROS 48000 S lyket al. 2008 3873.994 2.5(0.4) ... ESO FEROS 48000 Weselak et al. 2008 3873.994 3.11(0.23) ... g ESO VLT UVES 80000 Kre lowski et al.2012 3873.994 3.11(0.23) 1.502(0.068) h ESO VLT UVES 80000 Kre lowski et al.2012
82 3875.758 ≤ 1.0 ≤ 0.7 ESO CAT CES 60000 Gredel et al. 1991 3875.758 1.34(0.59) ... ESO CAT CES 120000 Palazzi et al. 1992 3875.758 1.3(0.2) ... ESO FEROS 48000 S lyket al. 2008 3875.758 1.0(0.2) ... ESO FEROS 48000 Weselak et al. 2008 3875.758 1.28(0.12) ... g ESO VLT UVES 80000 Kre lowski et al.2012 T01 2.2 ... ESO CAT CES 60000 Gredel et al. 1991 T01 2.77(0.33) ... ESO CAT CES 120000 Palazzi et al. 1992 T01 2.97(0.09) ... g ESO VLT UVES 80000 Kre lowski et al.2012 T01 2.93(0.07) ... h ESO VLT UVES 80000 Kre lowski et al.2012 Table 3.7: —Continued
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD 151932 CH+ 4232.548 12.5(0.1) 15.1 13.27(1.12) 15.36(1.3) ESO VLT UVES 80000 Weselak et al. 2009 A 4232.548 13.39(1.20) 15.13(3.38) ESO VLT UVES 80000 Weselak et al. 2009 B 4232.548 13.04(0.26) 15.09(0.30) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 4232.548 ≤ 20 15.01(0.34) k ESO VLT UVES 80000 Weselak et al. 2014 B 3957.689 6.93(1.60) ... ESO VLT UVES 80000 Weselak et al. 2009 B 3957.689 7.12(0.27) 15.02(0.57) ESO VLT m UVES m 80000 m Weselak et al. 2014 A CH 4300.313 23.0(0.1) 30.9 19.66(0.4) 51.9 ... m ... m ... m S lyket al. 2008 4300.313 ... 26.76(3.42) ESO VLT UVES 80000 Weselak et al. 2009 B 4300.313 ≤ 20 31.61(0.87) j ESO VLT UVES 80000 Weselak et al. 2014 B 4300.313 24.41(0.40) 31.25(0.61) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 3886.410 11.10(0.4) ...... m ... m ... m S lyket al. 2008 3886.410 5.97(0.78) 26.76(3.42) ESO VLT UVES 80000 Weselak et al. 2009 B 3886.410 5.75(0.41) 26.76(3.42) ESO VLT UVES 80000 Weselak et al. 2009 C m m m
83 3886.410 6.79(0.30) 16.29(0.72) ESO VLT UVES 80000 Weselak et al. 2014 A 3890.217 3.60(0.45) 26.76(3.42) ESO VLT UVES 80000 Weselak et al. 2009 B 3890.217 4.01(0.32) 26.76(3.42) ESO VLT UVES 80000 Weselak et al. 2009 C 3890.217 4.15(0.35) 15.00(1.24) ESO VLT m UVES m 80000 m Weselak et al. 2014 A CN 3874.602 7.7(0.1) 2.6 7.21(0.3) 2.87 ... m ... m ... m S lyket al. 2008 3874.602 7.20(0.40) 2.43 (0.07) ESO VLT UVES 80000 Weselak et al. 2009 B 3874.602 7.52(0.29) 1.773(0.054) h ESO VLT UVES 80000 Kre lowski et al.2012 3874.602 7.52(0.29) 1.685(0.064) g ESO VLT UVES 80000 Kre lowski et al.2012 3873.994 3.1(0.3) ...... m ... m m S lyket al. 2008 3873.994 2.56(0.20) ... ESO VLT UVES 80000 Weselak et al. 2009 B 3873.994 2.92(0.36) ... g ESO VLT UVES 80000 Kre lowski et al.2012 3873.994 2.92(0.36) 0.9950(0.076) h ESO VLT UVES 80000 Kre lowski et al.2012 3875.758 1.44(0.20) ... g ESO VLT UVES 80000 Kre lowski et al.2012 T01 3.27(0.26) ... g ESO VLT UVES 80000 Kre lowski et al.2012 T01 3.25(0.16) ... h ESO VLT UVES 80000 Kre lowski et al.2012 ... HD168076 CH+ 4232.548 ... 25.9 ... o 84.38(7.80) BOAO ... m ∼ 45000 n Weselak et al. 2014 B CH 4300.313 23.5(0.4) 29.5 ... 30(5) APO ARCES 38000 Thorburn et al. 2003 4300.313 ≤ 20 61.36(1.41) j BOAO ... m ∼ 45000 n Weselak et al. 2014 B CN 3874.602 ... 2.1 e ... 2(0.4) APO ARCES 38000 Thorburn et al. 2003 3874.602 5.5(0.2) 1.4 ...... Table 3.7: —Continued
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD 170740 CH+ 4232.548 13.78(0.08) 17.0 14.0(1.5) 18(2) ESO CAT CES 60000 Gredel et al. 1993 4232.548 13.97(0.06) ... ESO VLT UVES 75000 Casassus et al. 2005 4232.548 13.92(0.08) 16.11(0.09) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 4232.548 ≤ 20 16.80(0.11) k ESO VLT UVES 80000 Weselak et al. 2014 B 4232.548 13.63(0.26) 16.34(0.49) ESO VLT UVES 80000 Weselak et al. 2009 C 3957.689 8.0(1.6) 18(3) ESO CAT CES 60000 Gredel et al. 1993 3957.689 7.85(0.04) ... ESO VLT UVES 75000 Casassus et al. 2005 3957.689 7.99(0.07) 16.85(0.15) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 3957.689 8.01(0.23) 16.34(0.49) ESO VLT UVES 80000 Weselak et al. 2009 C CH 4300.313 15.32(0.04) 20.4 16.2(1.5) 24(3) ESO CAT CES 60000 Gredel et al. 1993 4300.313 ... 21(2) APO ARCES 38000 Thorburn et al. 2003 4300.313 16.01(0.05) 20.14(0.46) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 4300.313 ≤ 20 19.86(0.10) j ESO VLT UVES 80000 Weselak et al. 2014 B 3890.217 ≤2 ≤15 ESO CAT CES 60000 Gredel et al. 1993
84 3890.217 2.84(0.20) 22.02(0.57) ESO VLT UVES 80000 Weselak et al. 2010 3890.217 2.67(0.08) 9.57(0.20) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 3890.217 2.84(0.20) 22.02(1.39) ESO VLT UVES 80000 Weselak et al. 2009 C 3886.410 3.8(1.0) 19(5) ESO CAT CES 60000 Gredel et al. 1993 3886.410 5.10(0.29) 22.02(0.57) ESO VLT UVES 80000 Weselak et al. 2010 3886.410 4.29(0.06) 10.16(0.17) ESO VLT m UVES m 80000 m Weselak et al. 2014 A 3886.410 5.10(0.29) 22.02(1.39) ESO VLT UVES 80000 Weselak et al. 2009 C Table 3.7: —Continued
Present Literature Identifier Species Line Wλ N Wλ N Instrument Location Resolution Reference (A)˚ (mA)˚ (1012 cm−2) (mA)˚ (1012 cm−2) HD 170740 CN 3874.602 13.53(0.05) 6.2 14.0(0.4) 6.2(0.4) e ESO CAT CES 60000 Gredel et al. 1991 3874.602 14.0(0.4) 4.3(0.2) ESO CAT CES 60000 Gredel et al. 1991 3874.602 13.60(0.43) ... ESO CAT CES 120000 Palazzi et al. 1992 3874.602 ... 8.7 APO ARCES 38000 Thorburn et al. 2003 3874.602 13.47(0.05) 7.753 e ESO VLT UVES 80000 Kre lowski et al.2012 3874.602 13.47(0.05) 2.963(0.011) l g ESO VLT UVES 80000 Kre lowski et al.2012 3874.602 13.47(0.05) 5.519(0.036) h ESO VLT UVES 80000 Kre lowski et al.2012 3874.602 1.92(0.17) 8.55(2.38) ESO VLT UVES 80000 Weselak et al. 2009 C 3873.994 5.2(0.8) 1.9(0.3) ESO CAT CES 60000 Gredel et al. 1991 3873.994 6.68(0.61) ... ESO CAT CES 120000 Palazzi et al. 1992 3873.994 5.87(0.04) 2.234(0.026) lh ESO VLT UVES 80000 Kre lowski et al.2012 3873.994 0.66(0.14) ... ESO VLT UVES 80000 Weselak et al. 2009 C 3875.758 2.6(0.3) 1.9(0.3) ESO CAT CES 60000 Gredel et al. 1991 85 3875.758 3.82(0.60) ... ESO CAT CES 120000 Palazzi et al. 1992 3875.758 3.07(0.03) ... l ESO VLT UVES 80000 Kre lowski et al.2012 T01 2.8 ... ESO CAT CES 60000 Gredel et al. 1991 T01 3.49(0.22) ... ESO CAT CES 120000 Palazzi et al. 1992 T01 3.58(0.02) ... l ESO VLT UVES 80000 Kre lowski et al.2012 T01 2.72(0.02) ... h ESO VLT UVES 80000 Kre lowski et al.2012 3579.963 2.09(0.12) ... ESO VLT UVES 82000 Bhatt et al 2015 3579.963 1.88(01.0) ... g ESO VLT UVES 80000 Kre lowski et al.2012 3579.453 0.77(0.07) ... g ESO VLT UVES 80000 Kre lowski et al.2012 3580.937 0.037(0.04) ... g ESO VLT UVES 80000 Kre lowski et al.2012 T01 3.44(0.23) ... g ESO VLT UVES 80000 Kre lowski et al.2012 a3 sigma level. iReference noted that the given value is uncertain. bSame Dataset jBased on 4300 Awhen˚ W (4300 A)˚ ≤ 20 mA˚ cDerived from original data. kBased on 4232 Awhen˚ W (4232 A)˚ ≤ 20 mA˚ dDerived from original data, includes J = 1. lReference marked given line as saturated eIncludes J = 0 and J = 1 mNot stated f Assumed b = 1 nUnclear, listed as (30000, 45000, 90000) gBased on equivalent width with no saturation ef- o(4232 / 3957) (Doppler splitting) fects. hBased on fit with saturation effects taken into account. 3.2.2 HD 197770
HD 197770 has additional high resolution observations (≈ 3 km s−1 resolution) from the High Resolution Spectrograph (HRS) on the Hobby-Eberly Telescope (HET)
(Tull 1998) that were taken after my McDonald data . The resulting profile syntheses were provided by Dr. Adam Ritchey. There is good agreement between measurements of total column density for CH, CH+, and CN toward HD 197770, as shown in Table
3.8.
+ There is a small difference in the V lsr structure of CH which amounts to an offset of ' 2 km s−1. There is also a minor component seen in the HET data that is not seen in my McDonald data. Table 3.9 reveals that the component structure for
CH and CN are consistent.
The only remaining minor discrepancy comes from the b-value used for the CN
fitting. The current b-value for CN is only 0.8 km s−1, while the HET data have 1.8 km s−1. For CN (J = 1) lines, the total fit equivalent widths and column densities agree to within one sigma. For the R(0) line, there is still agreement at the one to two sigma level, but I have larger column density and a smaller equivalent width.
86 Table 3.8: Total Column Density and Wλ Comparison: HD 197770
Species Line/Band λ f -Value Wλ N Reference (A)˚ (mA)˚ (1012 cm−2) CH+ A-X (0,0) R(0) 4232.548 0.00545 2.1(0.7) 2.7(0.9) CURRENT A-X (0,0) R(0) 4232.548 0.00545 3.0(0.3) ... Cox et al. (2011) A-X (0,0) R(0) 4232.548 0.00545 2.74(0.15) 3.20(0.18) HET CH A-X (0,0) R(0) 4300.313 0.00506 26.0(1.7) 37.1(2.4) CURRENT A-X (0,0) R(0) 4300.313 0.00506 26.6(0.1) ... Cox et al. (2011) A-X (0,0) R(0) 4300.313 0.00506 26.36(0.10) 38.72(0.15) HET CN B-X (0,0) R(0) 3874.602 0.0337 15.6(0.6) 5.0(0.2) CURRENT B-X (0,0) R(0) 3874.602 0.0337 20.0(0.3) ... Cox et al. (2011) B-X (0,0) R(0) 3874.602 0.0337 16.64(0.43) 4.37(0.11) HET B-X (0,0) R(1) 3873.994 0.0225 5.4(0.6) 2.0(0.2) CURRENT B-X (0,0) R(1) 3873.994 0.0225 5.37(0.38) 2.02(0.14) HET B-X (0,0) P(1) 3875.758 0.0112 2.9(0.6) 2.0(0.4) CURRENT B-X (0,0) P(1) 3875.758 0.0112 2.90(0.38) 2.02(0.27) HET
87 Table 3.9: Component Column Density Comparison: HD 197770
CH+ CH CN V lsr N b V lsr N b V lsr N b Reference km s−1 1012cm−2 km s−1 km s−1 1012cm−2 km s−1 km s−1 1012cm−2 km s−1
88 -2.6 0.54(0.11) 3.2 -1.5 33.16(0.08) 1.4 -1.8 4.37(0.11) 1.8 HET 1.8 2.31(0.12) 3.5 1.8 5.56(0.12) 2.7 ...... HET 9.2 0.35(0.08) 2.2 ...... HET
-3.5 0.9(0.4) 1.0 ...... CURRENT -0.7 1.8(0.5) 1.5 -1.1 31.5(0.3) 1.1 -1.4 5.0(0.2) 0.8 CURRENT ...... 2.5 5.6(0.7) 3.0 ...... CURRENT 3.3 Fitting of the H2 Bands
All of our H2 spectra come from FUSE observations. The linelist for H2 was
1 based on that used by McCandliss . The majority of N (H2) is found in the low
J levels (0 and 1). These low J levels of H2 are very optically thick and most of the velocity component structure in those lines are so blended together that they
cannot be reliably determined by themselves. Section 3.3.2 discusses the H2 velocity component structure in more detail.
Figures 3-4, 3-5, and 3-6 show examples H2 spectra with different degrees of contamination from other sources and different degrees of optically thick absorption.
Each figure shows a spectrum covering four bands of H2: B-X (4-0), B-X (3-0), B-X (2- 0), and B-X (1-0). One of the four bands, B-X (1-0), typically has a higher degree of contamination and is only used as a final check by overlaying the fit from the other three bands. The H2 spectra for all other sight lines can be found in Appendix B. HD 108 in Figure 3-4 shows the least contamination of the three example spectra, while HD 170740 in Figure 3-5 shows considerably more stellar contamination, and
HD 192641 in Figure 3-6 has the most contamination, particularly around 1055 A.˚
The sample spectra show different degrees of H2 saturation. Fortunately, the wings of the optically thick low J levels of H2 still give a reliable measure of the column density. The relative populations of J = 0 and 1 provide a good estimate of the kinetic temperature, as collisions with H2 result in a thermal distribution between J = 0 and 1 levels (e.g., Savage et al. 1997):
" #!−1 hν g1N(0) Tkinetic = ln K. kb g0N(1)
1http://www.pha.jhu.edu/ stephan/H2ools/
89 The total statistical weight (g) is the product of the nuclear spin statistical weight, which is either 1 or 3 for para- or ortho- H2 respectively, and the rotational statistical weight (2J +1). The b-value for H2 is a combination of thermal and turbulent widths from the relative populations of J = 0 and 1 and CH, respectively. See Section 3.3.1 for a brief discussion.
3.3.1 b-Values
The b-values of H2 lines are a combination of thermal and turbulent broadening,
2 1/2 −1 b = (2kT/m + vturbulent) (km s ).
The turbulent broadening for each component of the H2 fits is based on the b-values of the individual component obtained from fitting the CH line at visible wavelengths for that line of sight. Because its mass is significantly greater than that of H2, the b-value of CH is primarily from turbulent broadening. The thermal broadening is calculated during the fit and is based on the relative populations of the J = 0 and 1 levels from H2.
3.3.2 Velocity Component Structure
Unlike the data from HST, the FUSE data does not have high enough resolution to fit the velocity component structure independently, and so the component struc- ture for H2 must come from another species. Federman (1982), Danks et al. (1984), Rachford et al. (2002), and Sheffer et al. (2008) all showed that N(CH) is linearly proportional to N(H2), linking these species. Thus, the velocity component structure of CH was the starting template used for the structure of each H2 line. This initial velocity component structure (from CH or other species) was kept constant during the fitting of H2 absorption.
90 3.3.3 Additional Constraints
Additional constraints were imposed to limit the range of possible solutions to
those that are allowed physically. One of the requirements imposed was setting limits
on the ranges of acceptable column density ratios between J levels. For J = 1, the
column density had to be between 700 and 25% of that for J = 0. For J = 2, the
column density had to be less than 105% of that for J = 1. For J ≥ 3, the column density for each level had to be between 3% and 50% of that for the J level below it. These ranges extend slightly beyond the range of physically reasonable excitation temperatures as well as the maximum and minimum observed values of the ratios found in previous work for similar clouds (Sheffer et al. 2008). Furthermore, these ranges are very helpful for identifying biases coming from stellar features only affecting certain J levels. When one of these limits is reached, it indicates some amount of additional care is required for an accurate fit; the spectra are then inspected for stellar features and other sources of error. In addition, the fits were all inspected to confirm the higher J levels being fitted were real features. There are several cases where upper limits are listed in the table of results (Table 3.10), meaning the feature was within the level of the noise.
91 3.3.4 H2 Results
Table 3.10 lists the column density results from the H2 fitting, providing individual J levels and totals for each sight line. There are two lines of sight with available
results from previous analyses with observations of the individual J levels, they appear
below our observed values in Table 3.10 and indicate the appropriate reference. The
values agree to within the errors for each sight line. The total error in NH2 is twice
the maximum difference in log space between the value of NH2 found by fitting the
normalized continuum and the values of NH2 found by fitting a scaled continuum that
has been either been raised or lowered by five percent. The NH2 error for each J level is calculated in the same way and uses 0.1 dex as the minimum error.
92 iue34 h gr hw h pcrmfrom spectrum the shows figure The 3-4: Figure
HD 108 rn h orbnso H of bands four the ering niae h tt h aa hw eea dots. as here shown data, the to fit the indicates 1.00
0.75
93 0.50 2 sdi h tig h oi line solid The fitting. the in used FUSE 0.25 oadH 0 cov- 108 HD toward
0.00
1040 1050 1060 1070 1080 1090 1100 iue35 h gr hw h pcrmfrom spectrum the shows figure The 3-5: Figure
HD 170740 te ore,eg,selrasrto n bopinfo unre- from absorption from and contamination species. absorption indicate lated stellar data e.g., the sources, and other fit the between ment oeigtefu ad fH of bands four the covering 1.00
0.75
94 0.50 2 sdi h tig h disagree- The fitting. the in used
FUSE 0.25 oadH 170740 HD toward
0.00
1040 1050 1060 1070 1080 1090 1100 iue36 h gr hw h pcrmfrom spectrum the shows figure The 3-6: Figure
HD 192641 hw ag ereo otmnto oprdt 3-4. to compared contamination of degree large H a of shows bands four the covering 1.00
0.75
95 0.50 2 sdi h tig hssgtline sight This fitting. the in used
FUSE 0.25 oadH 192641 HD toward
0.00
1040 1050 1060 1070 1080 1090 1100 Table 3.10: H2 Column Densities
Sight Line Log(NJ=0) Log(NJ=1) Log(NJ=2) Log(NJ=3) Log(NJ=4) Log(NJ=5) Log(NJ=6) Log(Total NH2 ) HD 108 20.52 (0.20) 20.27 (0.10) 18.70 (0.23) 18.17 (0.22) 17.61 (0.19) 16.68 (0.10) 16.15 (0.10) 20.72 (0.12) HD 5689 20.60 (0.10) 20.40 (0.10) 18.73 (0.11) 18.13 (0.19) 16.72 (0.19) ≤ 16.11 ≤ 14.66 20.82 (0.10) HD 12882 20.62 (0.10) 20.41 (0.17) 18.67 (0.16) 17.57 (0.10) 16.71 (0.10) 15.16 (0.12) 14.57 (0.10) 20.83 (0.10) HD 13841 20.41 (0.10) 20.14 (0.10) 18.49 (0.10) 17.46 (0.10) 16.47 (0.10) 15.43 (0.10) 14.88 (0.10) 20.60 (0.10) HD 13969 20.70 (0.10) 20.38 (0.14) 18.75 (0.14) 18.23 (0.14) 17.41 (0.10) 16.02 (0.10) 15.47 (0.10) 20.87 (0.10) HD 14053 20.77 (0.10) 20.51 (0.11) 18.02 (0.24) 17.50 (0.24) 16.97 (0.17) 15.59 (0.10) 15.02 (0.16) 20.96 (0.10) BD+56◦ 0501 20.58 (0.10) 20.25 (0.20) 18.67 (0.28) 18.13 (0.26) 17.56 (0.23) 16.74 (0.21) 16.21 (0.21) 20.75 (0.10) BD+56◦ 0563 20.38 (0.10) 20.22 (0.10) 19.01 (0.60) 18.48 (0.61) 17.60 (0.82) 15.88 (0.45) 15.34 (0.93) 20.62 (0.10) HD 14443 20.88 (0.10) 20.29 (0.10) 18.52 (0.10) 17.99 (0.10) 17.10 (0.10) 15.58 (0.10) 15.05 (0.10) 20.98 (0.10)
96 HD 14476 20.71 (0.13) 20.32 (0.18) 18.40 (0.10) 17.87 (0.10) 17.33 (0.10) 16.02 (0.10) 15.48 (0.10) 20.86 (0.14) BD+56◦ 0578 20.42 (0.15) 20.26 (0.12) 19.04 (0.21) 18.19 (0.71) 16.66 (1.00) 15.78 (1.10) ≤ 14.17 20.66 (0.14) HD 14947 20.65 (0.10) 20.25 (0.10) 18.52 (0.10) 17.63 (0.10) 16.58 (0.10) 15.17 (0.10) 14.51 (0.10) 20.80 (0.10) HD 15629 20.68 (0.10) 20.53 (0.16) 19.03 (0.10) 18.11 (0.10) 17.03 (0.10) 16.03 (0.10) 15.49 (0.10) 20.92 (0.10) HD 16691 20.69 (0.10) 20.10 (0.10) 19.38 (0.10) 18.70 (0.10) 17.85 (0.10) 16.03 (0.10) 15.05 (0.20) 20.81 (0.11) HD 17505 20.55 (0.10) 20.39 (0.10) 18.97 (0.10) 18.09 (0.15) 17.52 (0.74) 15.56 (0.45) 14.81 (0.10) 20.79 (0.10) HD 17520 20.78 (0.10) 20.18 (0.23) 19.17 (0.10) 18.64 (0.10) 18.12 (0.10) 17.34 (0.15) 16.81 (0.10) 20.89 (0.13) BD+60◦ 0586 20.62 (0.10) 20.01 (0.10) 19.04 (0.12) 18.30 (0.23) 17.76 (0.38) 17.24 (0.10) 16.71 (0.13) 20.72 (0.10) HD 19243 20.61 (0.12) 20.45 (0.10) 18.69 (0.47) 18.15 (0.47) 17.61 (0.48) 16.27 (0.88) ≤ 15.72 20.84 (0.10) HD 19820 20.52 (0.10) 20.26 (0.13) 18.08 (0.24) 17.55 (0.24) 16.95 (0.32) 15.32 (0.42) 14.47 (0.47) 20.71 (0.10) HD 25443 20.95 (0.19) 20.59 (0.17) 18.51 (0.82) 17.81 (0.51) 16.40 (0.40) 15.42 (0.41) ≤ 14.83 21.11 (0.18) HD 25638 20.80 (0.13) 20.59 (0.17) 18.03 (0.10) 17.71 (0.10) 16.35 (0.10) 15.85 (0.10) 14.31 (0.10) 21.01 (0.10) HD 45314 20.06 (0.25) 19.91 (0.20) 18.24 (0.11) 17.38 (0.10) 15.86 (0.10) 14.45 (0.10) 13.89 (0.14) 20.29 (0.25) HD 46223 20.39 (0.20) 19.95 (0.18) 18.72 (0.41) 17.79 (0.23) 16.85 (0.10) 15.73 (0.10) 14.61 (0.21) 20.53 (0.19) Table 3.10: —Continued
Sight Line Log(NJ=0) Log(NJ=1) Log(NJ=2) Log(NJ=3) Log(NJ=4) Log(NJ=5) Log(NJ=6) Log(Total NH2 ) HD 92964 20.64 (0.10) 20.42 (0.10) 17.38 (0.28) ≥ 16.86 ≥ 16.28 ≥ 15.71 ≥ 15.19 20.85 (0.10) HD 97253 19.80 (0.11) 19.59 (0.10) 18.03 (0.27) 16.72 (0.10) 15.34 (0.10) 13.85 (0.10) 13.30 (0.10) 20.02 (0.10) HD 108927 20.52 (0.10) 20.25 (0.10) 18.40 (0.10) 17.75 (0.43) 16.22 (0.10) 15.00 (0.10) 14.47 (0.10) 20.71 (0.10) HD 149404 20.62 (0.11) 20.05 (0.10) 18.29 (0.38) 17.21 (0.43) ≥ 16.03 ≥ 15.08 ≥ 14.50 20.73 (0.10) Jensen 2010 20.60 (0.03) 20.34 (0.05) 18.26 (0.06) 17.13 (0.06) 16.05 (0.04) 15.49 (0.04) 14.17 (0.04) ... HD 151932 20.59 (0.10) 20.31 (0.18) 18.50 (0.38) 17.87 (0.58) 17.15 (0.44) 16.07 (0.33) ≥ 15.54 20.77 (0.10) HD 168076 20.62 (0.12) 20.27 (0.10) 18.92 (0.10) 18.36 (0.12) 17.68 (0.78) 16.38 (0.19) 15.82 (0.51) 20.79 (0.10) HD 170740 20.62 (0.10) 20.40 (0.10) 18.90 (0.19) 17.38 (0.18) 16.13 (0.50) 15.60 (0.50) 15.06 (0.52) 20.83 (0.10) Jensen 2010 20.60 (0.05) 20.52 (0.11) 18.86 (0.02) 17.73 (0.03) 17.08 (0.05) 15.91 (0.17) 14.72 (0.26) ... HD 192303 20.60 (0.10) 20.39 (0.10) 18.82 (0.22) 18.28 (0.20) 17.76 (0.21) 17.02 (0.10) 16.50 (0.10) 20.81 (0.10)
97 HD 192641 20.52 (0.17) 20.37 (0.10) 19.29 (0.10) 18.26 (0.40) 17.71 (0.42) 16.74 (0.10) 16.20 (0.10) 20.77 (0.11) HD 193793 20.64 (0.10) 20.28 (0.17) 18.94 (0.80) 17.42 (0.70) 16.01 (0.50) ≥ 14.64 ≥ 14.08 20.80 (0.10) HD 197770 20.60 (0.10) 20.30 (0.10) 18.07 (0.10) 16.78 (0.10) 15.28 (0.10) ≥ 13.87 ≥ 13.29 20.78 (0.10) HD 214419 20.71 (0.50) 20.34 (0.52) 18.47 (0.60) 17.92 (0.75) ≥ 17.70 ≥ 16.51 ≥ 15.98 20.86 (0.66) HD 217086 21.00 (0.10) 20.59 (0.10) 18.78 (0.73) 17.25 (0.76) 16.05 (0.49) 14.54 (0.47) 13.99 (0.67) 21.15 (0.10) 3.4 Fitting of the CO Bands
The ISMOD routine iteratively fits all lines from each J level across all available
bands simultaneously. In most cases, the fit of each rotational line uses a fixed velocity
component structure based on the corresponding velocity component structure of CN
and CH seen at visible wavelengths with a fitted overall offset, as discussed in Section
3.4.1. Occasionally, several iterations of re-normalization and re-fitting were needed
to accurately create a normalization that reflected the true continuum level across
the different bands.
3.4.1 Velocity Structure
The initial velocity component structure of each line of 12CO (or 13CO) was based
on the structure of the deepest species observed for that line of sight, typically CN or
CH. CN is found to coexist with CO in a diffuse molecular cloud (Pan et al. 2004),
and it is the preferable tracer, but when only CN upper limits were available, the
velocity component structure from CH was used as the template.
The HST data, taken with the E140H grating, were of high enough resolution that the velocity component structure of CO could be independently confirmed to follow the structure of CN and CH. Unlike the CO data from HST, the FUSE data did not
have high enough resolution to fit the velocity component structure independently,
requiring the use of structure from other species. The high resolution HST spectrum
toward HD 108 (Figure 3-7) shows the detailed structure not observed in the medium
resolution HST observations of Figure 3-8 or in the lower resolution FUSE spectra
of Figure 3-9.
98 Figure 3-7: HD 108 CO spectra from HST (high resolution). Several of the A-X bands of CO are shown, A-X (7-0) to A-X (11-0). A-X (7-0) has the strongest f -value of the bands shown here and A-X (11-0) has the smallest f -value of the bands shown here. The locations of the R and Q branches are indicated on the figure for J levels 0 and 1. The spectra are centered about the rest velocity of the R(0) line in each band.
99 3.4.2 12CO Bands from HST
All the analyzed spectra from HST include the A-X (7-0), A-X (8-0), A-X (9-0),
A-X (10-0), and A-X (11-0) bands of CO. The medium resolution HST spectra also included the A-X (0-0), A-X (1-0), A-X (2-0), A-X (3-0), A-X (4-0), A-X (5-0), and
A-X (6-0) bands. The f -values for CO come from Chan et al. (1993) and Lambert et al. (1994). Table 3.11 displays a list of wavelengths and f -values. Figure 3-8, the medium resolution spectrum toward HD 108 from HST, illustrates the variation in band strength for the observed A-X bands. The weaker A-X bands, where the optical depth at line center is near one, are more sensitive to the total CO column density, while the stronger A-X bands allow the detection of weaker velocity components. The range of f -values between the A-X (0-0) to A-X (11-0) bands of CO leads to tighter constraints for the b-values for the components.
Table 3.11: HST 12CO Bands
12CO Band Band Head f -value (A)˚ (10−3) A-X (0-0) 1544.448 15.8 A-X (1-0) 1509.748 29.1 A-X (2-0) 1477.565 40.1 A-X (3-0) 1447.352 34.6 A-X (4-0) 1419.044 23.3 A-X (5-0) 1392.525 14.5 A-X (6-0) 1367.623 7.95 A-X (7-0) 1344.186 4.14 A-X (8-0) 1322.150 2.00 A-X (9-0) 1301.403 0.95 A-X (10-0) 1281.866 0.41 A-X (11-0) 1263.433 0.18
100 3.4.3 12CO Bands from FUSE
The FUSE spectra include many CO bands (only the main isotopic variant was seen in FUSE spectra and is denoted CO here); however, the majority of the CO bands have severe contamination from H2 absorption features. There are four CO bands with only small amounts of H2 absorption features: B-X (0-0), B-X (1-0), C- X (0-0), and E-X (0-0). The f -values for these CO bands come from the work of
Federman et al. (2001) and Sheffer et al. (2003). See Table 2.5 for the observational details and the f -values used. The B-X (1-0) band has the smallest f -value of the four bands shown. In several cases the B-X (1-0) band of CO is either too weak or cannot be separated from the stellar features, as in Figure 3-9. E-X (0-0) is situated on the wing of a strong H2 absorption line, but can be used in most cases after dividing the data by the fitted H2 spectrum. The C-X (0-0) band has a neighboring chlorine feature that makes this band dif-
ficult to fit on its own. Figures 3-9 and 3-10 show examples of the feature in the spectrum. Typically, the B-X (0-0), B-X (1-0) (if above the noise), and E-X (0-0) bands are fit together first. The result is overlaid on the C-X (0-0) band to look for agreement. The chlorine feature in C-X (0-0) can be analyzed after dividing the data by the fit of the CO band. All other CO observations from FUSE can be found in
Appendix C.
101 Table 3.12: HST 13CO Bands
13CO Band Band Head f -value (A)˚ (10−3) A-X (1-0) 1510.407 35.1 A-X (2-0) 1478.739 40.2 A-X (3-0) 1449.145 37.4 A-X (4-0) 1421.340 24.2 A-X (5-0) 1395.229 14.5 A-X (6-0) 1370.617 6.52 A-X (7-0) 1347.575 4.14
102 Figure 3-8: HD 108 CO spectra from HST (medium resolution). Several of the A-X bands of CO are shown, from A-X (1-0) to A-X (10-0). A-X (2-0) has the strongest f -value of the bands shown here and A-X (10-0) has the smallest f -value of the bands shown here. To the right of the A-X (1-0) band there is a large feature that is actually two overlapping bands of CO, d-X (5-0) of 12CO and A-X (1-0) of 13CO. See Figures 3-12 and 3-13 for details. The A-X (9-0) band also shows an unrelated absorption feature from P II near 100 km s−1. The locations of the R and Q branches for J levels 0 and 1 are only indicated below the A-X (10-0) band in the figure. The spectra are centered about the rest velocity of the R(0) line in each band.
103 Figure 3-9: This figure shows the CO spectra from FUSE toward HD 108. Three bands of CO are shown, B-X (1-0) is not shown because the weak features of the band cannot be differentiated from the noise and stellar features in the region. The B-X (1-0) band has the smallest f -values of the bands shown here and is often unable to be used. The locations of the different branches for J levels 0, 1, and 2 are indicated below the continuum of each band in the figure. The spectra are centered about the rest velocity of the R(0) line in each band. The C-X (0-0) band also shows an unrelated absorption feature from Cl around 50 km s−1 from the R(0) line. In each band, the solid line indicates the fit to the data, shown here as dots.
104 Figure 3-10: This figure shows the CO spectra from FUSE toward HD 25443. Four bands of CO are shown, including the weak B-X (1-0) band. The locations of the different branches for J levels 0, 1, and 2 are indicated below the continuum of each band in the figure. The spectra are centered about the rest velocity of the R(0) line in each band. The C-X (0-0) band also shows an unrelated absorption feature from Chlorine around 50 km s−1 from the R(0) line. The solid lines indicate the fit to the data, shown here as dots.
105 3.4.4 13CO Detections with HST
There are only a few detections of 13CO from the HST spectra. Table 3.12 displays
a list of wavelengths and f -values for this isotopologue. A few illustrative spectra are
provided here as examples; all other 13CO spectra are found in Appendix D. Figure
3-11 shows a clear detection of the A-X (7-0) band of 13CO toward HD 170740 and
a unrelated chlorine feature to the left. Figure D-3 shows a similar spectrum toward
HD 108927. Within the spectrum toward HD 108, containing the A-X (1-0) band
of CO (Figure 3-8), there is another strong feature to the right that is actually two
overlapping bands of CO, d-X (5-0) of 12CO and A-X (1-0) of 13CO. Figure 3-12 shows the fit from the other CO bands overlaid on the d-X (5-0) band of 12CO, and Figure
3-13 shows the residual 13CO spectrum. Unfortunately, this is the only available measure of 13CO toward HD 108 and because of the large residuals and uncertainty from overlapping bands of CO, it is only listed as an upper limit. Figures D-5 and D-6 show similar spectra toward HD 46223. Figure 3-14 shows the 13CO spectra toward
HD 46223 in all observed and detected bands.
3.4.5 12CO and 13CO Results
The combined results of the CO fitting for HST and FUSE data are listed in
Table 3.13: column densities of the individual J levels up to J = 3, the total column density for each line of sight, and the excitation temperatures for the observed levels.
There are only a few detections of 13CO from the HST spectra. The results of 13CO are included when the isotopologue is detected; in some cases only upper limits can be reliably determined.
106 Figure 3-11: The figure shows the spectrum toward HD 170740 from HST covering the A-X (7-0) band of 13CO. The locations of the R and Q branches for J levels 0 and 1 are indicated. The spectrum is centered about the rest velocity of the R(0) line. The A-X (7-0) band also shows an unrelated absorption feature from chlorine near -100 km s−1.
107 Figure 3-12: The figure shows the spectrum toward HD 108 from HST cov- ering the A-X (1-0) band of 13CO; unfortunately this coincides with the d-X (5-0) band of 12CO. To the left of the feature is the A-X (1-0) band from 12CO. The solid line shows the fit from the other CO bands overlaid on the d-X (5-0) band of 12CO. Figure 3-13 shows the residual 13CO spectrum.
108 Figure 3-13: The figure shows the spectrum toward HD 108 from HST cov- ering the A-X (1-0) band of 13CO with the overlapping d-X (5-0) band of 12CO removed.
109 Figure 3-14: The figure shows all HD 46223 spectra (medium resolution) from HST with detected 13CO. The locations of the branches are shown below the A-X (1-0) spectrum. The nearby 12CO bands in the A-X (1-0) and A-X (6-0) spectra are indicated above the continuum.
110 Table 3.13: CO Results
Sight Line NJ=0 NJ=1 NJ=2 NJ=3 Total NCO T(0−1) T(0−2) T(0−3) (1014cm−2) (1014cm−2) (1014cm−2) (1014cm−2) (1014cm−2) (K) (K) (K) HD 108 5.0(0.1) 3.8(0.1) 0.46(0.12) ... 9.3(0.2) 4.0(0.1) 4.2(0.3) ... HD 108 (13CO) ≤ 0.6 ≤ 1.2 ...... ≤ 1.8 ...... HD 5689 2.4(0.2) 1.5(0.2) ≤ 0.14 ... 3.9(0.3) 3.5(0.3) ...... HD 12882 12.0(0.6) 8.7(1.1) ≤ 0.60 ... 20.7(1.2) 3.9(0.4) ...... HD 13841 2.4(0.2) 2.8(0.4) ≤ 0.21 ... 5.2(0.5) 5.9(1.2) ...... HD 13969 8.3(1.5) 5.8(1.3) ≤ 0.65 ... 14.1(2.0) 3.8(0.8) ......
111 HD 14053 1.6(0.5) 1.6(0.5) ≤ 1.20 ... 3.2(0.7) 5.0(2.0) ...... BD+56◦ 0501 2.1(0.2) 1.6(0.3) ...... 3.7(0.4) 4.0(0.6) ...... BD+56◦ 0563 1.1(0.2) 1.3(0.4) ...... 2.4(0.5) 5.9(2.3) ...... HD 14443 1.1(0.3) 1.2(0.3) ≤ 1.29 ... 2.3(0.4) 5.5(2.1) ...... HD 14476 3.7(1.0) 4.5(1.4) ...... 8.2(1.7) 6.1(2.8) ...... BD+56◦ 0578 1.1(0.1) 1.4(0.2) ...... 2.5(0.2) 6.4(1.2) ...... HD 14947 21.0(8.9) 9.8(4.1) ...... 30.8(9.8) 3.0(1.0) ...... HD 15629 7.4(2.3) 7.3(2.3) ≤ 6.15 ... 14.7(3.3) 5.0(2.0) ...... HD 16691 15.0(1.5) 7.4(1.4) ≤ 6.60 ... 22.4(2.0) 3.1(0.4) ...... HD 17505 12.0(1.1) 6.6(1.1) ≤ 1.33 ... 18.6(1.5) 3.3(0.4) ...... HD 17520 2.3(0.2) 2.6(0.6) ≤ 0.55 ... 4.9(0.6) 5.7(1.5) ...... BD+60◦ 0586 3.7(0.4) 3.5(0.4) 0.59(0.13) ... 7.8(0.6) 4.8(0.6) 4.8(0.3) ... HD 19243 2.1(0.2) 2.5(0.3) ...... 4.6(0.4) 6.0(1.0) ...... Table 3.13: —Continued
Sight Line NJ=0 NJ=1 NJ=2 NJ=3 Total NCO T(0−1) T(0−2) T(0−3) (1014cm−2) (1014cm−2) (1014cm−2) (1014cm−2) (1014cm−2) (K) (K) (K) HD 19820 18.0(2.9) 12.0(3.1) ...... 30.0(4.2) 3.7(0.7) ...... HD 25443 7.2(0.3) 2.4(0.4) ≤ 0.39 ... 9.6(0.5) 2.5(0.2) ...... HD 25638 15.0(2.9) 8.0(2.4) ≤ 1.73 ... 23.0(3.7) 3.2(0.6) ...... HD 45314 4.1(0.1) 3.1(0.3) 0.32(0.10) ... 7.5(0.3) 4.0(0.3) 4.0(0.1) ... HD 46223 3.4(0.1) 1.8(0.1) 0.21(0.01) ... 5.4(0.1) 3.2(0.2) 3.8(0.1) ... HD 46223 (13CO) 0.09(0.01) 0.06(0.02) ...... 0.15(0.02) 3.7(0.9) ...... HD 92964 6.2(0.9) 7.9(1.1) ≤ 3.27 ... 14.1(1.4) 6.5(1.5) ...... HD 97253 2.3(0.1) 2.0(0.1) ≤ 0.30 ... 4.3(0.1) 4.5(0.2) ......
112 HD 108927 9.0(0.4) 9.3(0.5) 1.10(0.19) ... 19.4(0.6) 5.2(0.3) 4.5(0.2) ... HD 108927 (13CO) 0.24(0.08) 0.25(0.17)a ...... 0.49(0.19)a 5.2(3.8)a ...... HD 149404 5.6(0.4) 5.7(0.5) 1.10(0.14) ... 12.4(0.6) 5.1(0.5) 5.1(0.2) ... HD 151932 4.7(0.5) 4.7(0.6) 0.85(0.25) ... 10.3(0.9) 5.0(0.8) 5.0(0.5) ... HD 168076 2.6(0.3) 2.3(0.3) ≤ 0.38 ... 4.9(0.4) 4.5(0.7) ...... HD 170740 41.0(1.8) 58.0(2.8) 13.00(1.66) 1.6(0.3) 113.6(3.7) 7.4(0.6) 6.0(0.3) 6.4(0.3) HD 170740 (13CO) 0.59(0.14) 0.89(0.27) ...... 1.48(0.30) 8.0(4.4) ...... HD 192303 1.1(0.2) 1.4(0.3) ≤ 1.50 ... 2.5(0.4) 6.4(2.1) ...... HD 192641 0.6(0.2) 0.8(0.3) ≤ 0.62 ... 1.5(0.3) 6.7(3.8) ...... HD 193793 1.1(0.1) 1.0(0.1) 0.13(0.04) ... 2.2(0.2) 4.4(0.5) 4.4(0.5) ... HD 197770 14.0(1.2) 8.7(2.0) 0.61(0.16) ... 23.3(2.3) 3.5(0.5) 3.5(0.2) ... HD 214419 41.0(8.2) 19.0(4.8) ≤ 1.21 ... 60.0(9.5) 3.0(0.5) ...... HD 217086 6.4(1.4) 4.9(1.1) ...... 11.3(1.8) 4.0(0.9) ...... aThe J = 1 column density for 13CO is inferred from the fitted J = 0 column density and the 12CO excitation temperature. 3.5 Fitting of Additional Species
There are several additional species that are only available for a small subsample of sight lines. There are also additional sight lines and species observed as part of the
GOT C+ project.
3.5.1 C I
The sight lines observed with HST included observations of several C I multiplets.
Unlike most of the fits with other species, the fraction and b-value for each component in each fine-structure level were allowed to vary. The weaker, optically thin multiplets are more sensitive to the total column density, while the stronger multiplets are more sensitive to the presence of minor components. The final fit for each sight line is based on the simultaneous fit of the lines from all detectable fine-structure levels across the observed multiplets. Table 3.14 lists the total column densities for each fine-structure level, and Table 3.15 gives the Vlsr, fraction, and b-value for each component.
Table 3.14: C I Column Densities
Star N J=0 N J=1 N J=2 N T otal (1013 cm−2) (1013 cm−2) (1013 cm−2) (1013 cm−2) HD 108 545(8) 98.2(0.3) 23.0(0.2) 666(9) HD 13841 248(3) 21.4(0.4) 10.6(1.2) 280(7) HD 25443 490(3) 36.6(0.3) 6.6(0.1) 533(3) HD 46223 907(8) 53.9(0.3) 9.4(0.1) 970(8) HD 108927 492(5) 54.2(0.4) 10.9(0.2) 557(5) HD 170740 421(3) 49.6(0.3) 24.0(0.3) 495(11)
113 Table 3.15: C I Component Structure
Star Vlsr f b-value J = 0 J = 1 J = 2 J = 0 J = 1 J = 2 (km s−1) (km s−1) (km s−1) (km s−1) HD 108 -16.8 0.08 0.06 0.05 3.1 3.5 3.2 -9.7 0.11 0.04 0.04 2.4 1.4 1.7 -5.5 0.39 0.49 0.39 1.3 1.2 1.4 -1.7 0.10 0.03 0.05 3.0 1.6 0.6 2.4 0.10 0.03 0.05 3.5 1.7 2.3 6.6 0.22 0.34 0.42 0.7 0.9 1.0
HD 13841 -41.2 0.12 0.12 .001 3.5 1.3 0.5 -15.6 0.09 0.07 0.02 2.7 1.8 2.0 -10.5 0.42 0.56 0.95 0.8 0.7 0.3 -2.9 0.37 0.25 0.03 3.5 3.5 0.3
HD 25443 -10.9 0.28 0.40 0.42 3.5 1.9 2.1 -7.2 0.16 0.10 0.07 0.6 0.4 0.3 1.3 0.23 0.19 0.20 1.6 1.5 1.0 6.4 0.33 0.31 0.31 1.1 0.8 0.6
HD 46223 -5.8 0.01 0.03 0.06 1.1 0.9 3.5 1.4 0.12 0.14 0.11 0.6 0.5 0.7 6.6 0.06 0.19 0.17 2.8 1.0 2.0 8.7 0.67 0.32 0.31 0.9 1.0 1.8 14.4 0.14 0.33 0.34 1.9 1.8 3.5
HD 108927 18.2 0.99 0.97 0.98 1.3 1.0 1.1 21.2 0.01 0.03 0.02 2.8 0.3 3.5
HD 170740 -12.5 0.23 0.23 0.08 2.4 0.9 0.6 -9.2 0.77 0.77 0.92 0.7 0.9 0.7
114 3.5.2 C2
The wavelengths of the R(0) line in each C2 band are given in Table 3.16 along with the f -values and references. Sonnentrucker et al. (2007) adapted the wavelengths in both bands from van Dishoeck & de Zeeuw (1984). The A-X (2-0) f -value in
Sonnentrucker et al. (2007) came from averaging the values from Erman & Iwamae
(1995) and Langhoff et al. (1990). Sonnentrucker et al. (2007) provides a listing of wavelengths and relative f -values for each line, given in relation to the R(0) f -value provided here.
The results of the C2 column density fits are given in Tables 3.17 and 3.18. Figure
3-16 shows the C2 A-X (2-0) band toward HD 108927 from VLT/UVES observations
(λ/∆λ ≈ 88,000) and Figure 3-15 shows the C2 A-X (3-0) toward HD 197770 from the HRS/HET observations (λ/∆λ ≈ 100,000); both spectra were reduced by Dr. Adam
Ritchey. The residuals (data - fit) are shown as the top spectrum shifted up by 0.05.
The middle spectrum provides the fit overlaid on the data and the bottom spectrum shows just the fit, shifted down by 0.05. This is done to aid in differentiating the weaker C2 features. For HD 197770, the R(10) line of C2 was excluded from the fit because it falls on a small residual emission feature from a cosmic ray.
Table 3.19 provides the total C2 column density for several sight lines observed in previous studies. The C2 features toward HD 197770 were also observed at high resolution (80,000) by Hanson et al. (1994), but their column density is slightly more than twice our observed value. Unfortunately, Hanson et al. (1994) do not reference the f -value they used, which could be the source of the discrepancy.
115 Table 3.16: Bands of C2
C2 Band Wavelength f -value Reference (A)˚ (10−3) A-X (2-0) 8757.686 1.40 Sonnentrucker et al. (2007) A-X (3-0) 7719.329 0.65 Sonnentrucker et al. (2007)
Table 3.17: C2 Structure
Star Vlsr fraction b (km s−1) (km s−1) HD 108927 1.5 1.00 0.8 HD 197770 -3.8 0.10 2.5 -1.4 0.90 0.7
Table 3.18: C2 Column Density
a a a a a a Star NJ=0 NJ=2 NJ=4 NJ=6 NJ=8 Ntotal HD 108927 0.6 (0.4) 4.2 (0.7) 3.9 (0.7) 2.4 (0.7) 2.7 (0.7) 13.7 (3.4) HD 197770 0.7 (0.5) 7.1 (0.5) 3.7 (0.5) 4.3 (0.5) 1.5 (0.7) 17.3 (3.9) aColumn densities are in units of 1012cm−2.
Table 3.19: Other C2 Results
a Sight Line N(C2) Reference HD 149404 ≤ 20 van Dishoeck & de Zeeuw (1984) HD 197770 32 Hanson et al. (1994) HD 151932 26 (2) Ka´zmierczaket al. (2010) HD 168076 48 (13) Thorburn et al. (2003) HD 170740 18 (2) Ka´zmierczaket al. (2010) 24 (14) Thorburn et al. (2003) aN is given in units of 1012 cm−2.
116 iue31:TefiuesosteC the shows figure The 3-15: Figure
1.05 n h otmsetu sjs h t hfe onb 0.05. C by the down differentiating shifted in data fit, aid the to the on by done just overlaid is up is fit This shifted spectrum the bottom spectrum, shows the top spectrum and the middle are The fit) 0.05. - (data residuals The
1.00 117 2 A-X 30 adtwr D197770. HD toward band (3-0)
0.95 2 features. HD197770 HD197770
0.90 7710 7720 7730 7740 iue31:Asrto pcr fC of spectra Absorption 3-16: Figure
1.05 h idesetu hw h tt h aa e Figure See data. the to fit details. additional the for shown 3-15 spectrum middle The
1.00 118 2 A-X 20 adtwr D108927. HD toward band (2-0)
0.95
HD108927
0.90 8750 8760 8770 3.5.3 Species for the GOT C+ Project
The data for the GOT C+ pointings include the archival emission spectra of C+,
CO and its isotopologues, and, H I, from the Herschel Space Observatory, the ATNF
Mopra Telescope, and the VLA Galactic Plane Survey (Stil et al. 2006). The data were fitted with 1-3 Gaussians on top of a linear baseline using the MPFIT package from Markwardt (2009). The initial fits for the Gaussians were determined by eye.
This fitting was performed by Dr. Nicolas Flagey.
Each sight line from this project has a stacked plot showing the spectra of the species that were observed in absorption toward the star and in emission for the closest GOT C+ pointing. Only Figure 3-17 toward HD 165783 is shown, while all remaining figures can be found in Appendix E. Species not shown were not observed for that pointing, with the exception of C18O, which is not shown due to the low correspondence with absorption components. The Vlsr of components detected at a three sigma level are indicated for each species as a red tick above the spectra.
Our results for both the absorption and emission lines in our study are combined into Tables 3.20 and 3.21. These tables show the Vlsr for components with emission features and any corresponding absorption features for the inner or outer Galaxy. The component structure seen in absorption at visible wavelengths from Ca II, Ca I,K I,
CH, CH+, and CN is combined with emission from H I, CO and its isotopologues, and
C+ from the GOT C+ survey. The correspondence between components in emission and absorption help create a more unified picture of diffuse atomic and molecular gas in the interstellar medium.
The following points were considered when seeking matches between emission and absorption component velocities. First, diffuse atomic and molecular gas, with kinetic temperatures of 50 to 80 K, have thermal widths of about 1.0 km s−1. The extensive survey conducted by Pan et al. (2005) revealed typical b-values of 1 to 2
119 km s−1 for absorption lines, depending on species. Since the species in our survey
from the McDonald observations have atomic masses greater than 20 amu, thermal
broadening makes a negligible contribution to the b-value, which mainly arises from
turbulent motion after removing the instrumental width of 2.2 km s−1. Second, the
light path for the hollow cathode lamp used in wavelength calibration differs slightly
from that of the stellar radiation. This difference causes 0.5 to 1.0 km s−1 offsets from an absolute velocity scale. Third, most of the background stars in our survey have spectroscopic distances within 500 pc of the Sun, with only one significantly beyond
1000 pc. On Galactic scales, these stars are considered within the solar neighborhood.
The majority of the gas along the lines of sight have velocities in the Local Standard of Rest less than 10 km s−1. Dynamical phenomena like stellar winds and expanding supernova remnants create components with larger radial velocities, some of which are in excess of 100 km s−1. An example involves directions that probe the supernova remnant IC 443 (Welsh & Sallmen 2003; Hirschauer et al. 2009). This leads to complications in associating components seen in absorption and emission in the inner
Galaxy where large velocities may arise from dynamical phenomena, especially in atomic gas, or Galactic rotation. In light of these considerations, correspondences were noted when velocities associated with line centers agreed within the full width at half maximum of the lines. Weighted averages based on column density were performed for some species when comparisons were made with the broadest lines
(usually H I).
When comparing observations in absorption and in emission, there will frequently be components that appear shifted, blended, or even missing entirely. Over half (56%) of the emission components are not associated with an absorption component, and over half (53%) of the absorption components are not associated with an emission component. The majority of the unassociated components are only seen in species associated with low density gas. Over half (52%) of the unassociated emission com-
120 ponents are only seen in H I and 80% of the unassociated absorption components are
only seen in Ca II.
There are several sources of these inconsistencies. One is the fact that absorption lines only sample gas along the line of sight between the background source and the observer, while the emission lines sample all the gas. This is the reason components observed in emission at high Vlsr are sometimes absent in absorption. The Vlsr struc- ture for two nearby absorption sight lines can exhibit similar differences when the background sources are at different distances.
Secondly, at low densities almost all the gas seen in absorption is in the lowest states and the maximum column density is probed while emission lines typically come from weakly excited states in diffuse gas and are sensitive to density. This difference results in the absorption measurements being more sensitive to detecting the gas than the emission lines, but is also restricted to lines of sight with background sources while emission lines can be mapped. This is one of the sources for additional components seen in absorption that are not seen in emission. Over half (80%) of the unassociated absorption components are only seen in Ca II, which indicates low density gas.
Another source of inconsistency is from the blending of lines. Sometimes the rel- atively broad emission lines result in a component that corresponds to a blend of two absorption components with slightly different velocities. Whenever two components had to be combined to form a corresponding component, it is indicated by the use of the superscript ‘a’ on the Vlsr in Tables 3.20 and 3.21, where an averaged velocity (weighted by column density) is quoted for the McDonald data.
There are occasional additional components seen in emission that are not seen as absorption components and fall within the overlapping Vlsr ranges. The most likely explanation is from variations in the small scale structure between the differ-
ent locations of the emission and absorption measurements. In general, whenever an
observed species is associated with one or more of the unobserved species, those unob-
121 served species are likely present at varying levels between the emission and absorption
pointings. However, at the location of the measurements, the levels are too low to
be confidently measured. Table 4.14 shows a summarized component-by-component
interpretation for a pointing with two nearby sight lines. The majority of compo-
nents for the two sight lines are the same, but there are also clear differences. The
−1 component with a Vlsr of 7 km s is seen in emission toward the GOT C+ pointing and in absorption toward HD 165783 but not toward the other nearby sight line, HD
165918. This small scale structure can also cause components to be slightly shifted in velocity.
122 Figure 3-17: (Top) Absorption spectra in Vlsr toward HD 165783. (Bottom) Emission spectra in Vlsr from the closest pointing, G010.4+0.0. The individual components are shown as dotted lines in each species and indicated by red ticks above the fit. A dashed line for each individual component is also shown indicating the con- tribution to the overall fit; this is only shown for this sight line.
123 Table 3.20: Velocities of Components with Corresponding Emission Features in the Inner Galaxy
−1 Species Vlsr (km s ) G010.4+0.0 C+ -38.9 ...... 5.8 ...... 25.0 ... 45.3 65.0 ... 73.7 ...... H I ... -25.8 -19.5 -10.0 1.2 ...... 11.8 20.9 ... 35.4 45.3 65.7 ... 75.1 90.7 122.4 ... 150.9 12CO ...... -11.0 ...... 6.8 ... 20.4 26.7 34.0 ...... 69.5 ...... 122.9 145.7 153.2 13CO ...... -11.0 ...... 6.3 ... 20.5 23.4 34.3 ...... 69.0 ...... C18O ...... 68.8 ...... 124 HD 165783 Ca II ...... -21.5 -10.2a 1.1 4.9 7.7 10.7a 20.4 26.7a 34.8 45.2 ...... Ca I ...... -10.5 ... 6.0 ...... CH+ ...... -9.1 ... 5.0 6.9 ...... CH ...... 4.1 7.3 ...... HD 165918 Ca II ...... -21.2 -10.3a 0.0 4.2 ... 10.7 ...... CH+ ...... -7.7 ... 5.1 ...... CH ...... 4.9 ...... Table 3.20: —Continued
−1 Species Vlsr (km s ) G014.8-1.0 C+ ...... 20.9 ...... H I -24.1 -8.3 5.3 15.0 ... 23.5 ... 37.4 ... 54.9 70.0 112.6 12CO -26.1 ...... 19.4 ... 30.0 39.1 46.3 55.9 75.6 ... 13CO ...... 19.3 ... 29.7 ...... 74.8 ... C18O ...... 29.9a ......
HD 168607 Ca II -23.0 -9.8 6.7 15.0 18.7 23.3a 30.7 37.9a 46.2 52.6 ...... 125 CH+ ...... 5.5 15.2 ... 22.7a ...... CH ...... 7.0 14.5 18.3 22.9 ...... CN ...... 18.5 ......
G015.7+1.0 C+ -79.5 ...... 18.1 28.5 ...... H I ... -41.3 -19.6 -5.0 5.0 ... 18.7 30.2 ... 45.0 61.2 95.0 140.9 12CO ...... 10.0 19.1 29.2 38.2 48.4 ......
HD 167498 Ca II ...... -6.0 4.1 ...... CH+ ...... -6.5 4.4 ...... CH ...... 5.1 ...... Table 3.20: —Continued
−1 Species Vlsr (km s ) HD 167812 Ca II ...... 4.1 ...... CH ...... 5.4 ......
G020.0+0.0 C+ ...... 28.2 ... 37.1 46.4 55.9 67.8 ...... 114.5 ...
126 H I -28.8 -12.6 -0.8 ... 8.8 20.4 ... 34.0 ... 44.3 54.9 65.0 76.8 80.9 94.3 114.2 122.6 12CO ...... 6.6 ...... 27.8 ...... 45.4 ... 65.0 72.1 80.3 ... 118.8 ... 13CO ...... 7.0 ...... 27.8 ...... 45.4 ... 65.0 71.3 ...... C18O ...... 65.0 71.2 ......
HD 169754 Ca II ... -11.6 1.1a 5.5 ... 19.0 26.4 32.3 38.2 44.0a ... 63.5 ...... Ca I ...... 2.5 ...... 22.9 ...... CH+ ...... -0.2a 4.6 ... 21.4 ...... CH ...... 1.5a 5.8 ...... CN ...... 0.5a 5.5 ...... Table 3.20: —Continued
−1 Species Vlsr (km s ) G032.6+0.0 C+ ...... 16.0 26.6 ... 52.9 ... 74.4 ... 101.1a H I -48.0 -37.2 -19.3 -0.5 ...... 14.4 25.3 35.0 52.4 ... 76.9 91.2 102.4 12CO ...... 5.8 9.7 ...... 71.6 77.4 93.5 101.5 13CO ...... 9.5 12.1 ...... 74.4a 93.3 101.6 C18O ...... 74.1 ...... 127 HD 174509 Ca II ...... 0.1 2.8 7.8 12.2 15.7 ...... Ca I ...... 1.0 3.2 8.9 11.6 14.9 ...... K I ...... 1.0 3.8 8.4 11.5 ...... CH+ ...... 0.3 ... 7.6 10.5 ...... CH ...... 0.7 3.5 8.4 12.0 ...... CN ...... 3.2 7.4 11.3 ......
aOccasionally two nearby components need to be combined to create agreement with the features in another species. This is done by weighting the Vlsr of the components by their relative column densities and taking an average. Table 3.21: Velocities of Components with Corresponding Emission Features in the Outer Galaxy
−1 Species Vlsr (km s ) G091.7+1.0 C+ ... -100.7 ...... H I -111.9 ... -93.9 -82.2 -70.5 -54.0 -43.5 -32.6 -10.1 0.1 5.0 128 BD+49◦ 3482 Ca II ...... 1.6 5.4a K I ...... 1.2 ... CH+ ...... 1.5 3.7
BD+49◦ 3484 Ca II ...... 3.5 CH+ ...... 2.8 Table 3.21: —Continued
−1 Species Vlsr (km s ) G109.8+0.0 C+ ...... -50.4 ...... H I -92.8 -75.0 -59.4 -47.0 -39.0 -10.0 1.8 3.7
HD 240179 Ca II ...... -10.3 0.2 4.0 Ca I ...... -9.1 ... 3.9 K I ...... 3.5 CH+ ...... -9.9 ......
129 CH ...... 0.3 ...
HD 240183 Ca II ...... -10.4a 2.1 ... G207.2-1.0 C+ ...... 13.0 ...... H I 3.2 9.2 ... 19.7 25.0 38.3 12CO ... 8.7 ......
HD 47073 Ca II 3.8 ...... HD 260737 Ca II 1.8 9.4a ...... K I ... 7.2 ...... Table 3.21: —Continued
−1 Species Vlsr (km s ) G225.3+0.0 C+ 17.8 51.9 12CO 17.4 ...
HD 55469 Ca II ...... K I ...... CH+ ......
130 HD 55981 Ca II 14.8b ... K I 14.9b ... CH+ 14.5b ...
aOccasionally two nearby components need to be combined to create agreement with the features in another species. This is done by weighting the Vlsr of the components by their relative column densities and taking an average. bThis component is likely associated with the emission component through additional correspondence with data from previous work. (See text for further details.) Chapter 4
Analysis
This chapter explores different techniques and measures of density. In Section
4.1, the trends in the relationship between CO versus H2 column density are used as a tool to explore changes in the chemistry and to make comparisons in density for
the observed sight lines. In Section 4.2, excitation temperatures are used to provide
additional measures of density. Section 4.3 describes how steady state chemistry can
be used to infer the density. In Section 4.4, we apply several different density deter-
minations to analyze our h and χ subsample and resolve an apparent inconsistency
between different density measures. In Section 4.5, measurements of 13CO are used to
explore differences in physical conditions and processes in the 13CO containing gas.
Section 4.6 looks at kinematic associations with the species observed in the GOT
C+ survey to deduce correspondences between density tracers seen in emission ver-
sus absorption. All methods yielded average densities for sight lines. The analyses
utilizing C I excitation and CN chemistry were also performed for individual velocity
components.
4.1 Column Density Trends
The first indication of a relationship between CO and H2 column densities was made by Federman et al. (1980). The change in the slope of the power law between 131 N (CO) and N (H2), as seen in Sheffer et al. (2008), demonstrate a change in the CO photochemistry as CO starts to self shield and the density increases causing a different chemical chain to begin to dominate. Self shielding occurs on a line by line basis, when the photoexcitation transition becomes optically thick, with stronger self shielding occurring for stronger lines and higher J level populations. Sheffer et al.
(2008) show that at low molecular hydrogen column densities (log[N (H2)] less then
20.4), there is a linear relation between N (CO) and N (H2) with a slope of 1.46(0.23).
When the CO photochemistry changes, the relation between N (CO) and N (H2) also changes. The slope above the break is 3.07(0.73) (Sheffer et al. 2008).
The full chemical network used in the analysis by Sheffer et al. (2008) includes over
1200 reactions and is available from the Cloudy website.1 We outline the dominant
production and destruction pathways for CO here. At low N (CO), the CH++O
reaction begins the primary production route of CO,
CH+ + O → CO+ + H.
This is followed by conversion of CO+ to CO by
+ + CO + H2 → HCO + H, and HCO+ + e → CO + H, or
CO+ + H → CO + H+.
At N (CO) ' 1014 cm−2, the CH+ chemistry is less important due to the enhanced
destruction of CH+ at the higher densities in the gas and the C++OH production
route for CO begins to dominate:
C+ + OH → CO+ + H.
1http://www.nublado.org 132 This is followed by conversion of CO+ to CO by the same set of reactions given above.
Photodissociation is the primary destruction pathway for CO. As the column
density of CO increases, CO starts to self shield and quickly decreases the impact of
photodissociation, while the OH pathway is still effectively forming CO. This produces
the kink in the trend between N (CO) and N (H2). Figure 4-1 shows an updated version of the figure in Sheffer et al. (2008) with the additional data points from our
measurements. Sheffer et al. (2008) find that this relation can be extrapolated into
the center of the dark clouds distribution from Federman et al. (1990) with the N (H2)
values not directly observed but inferred from their visual extinction, AV . Figure 4-1 shows that our current data set, on the whole, is evenly distributed
about the average trend lines for the different regimes. On average, our current data
set also tends to be found at larger N (H2) than the Sheffer et al. (2008) data set. As a result, we do expect to see differences in several measures of the physical conditions
in the gas, such as the average excitation temperature and the species able to be
observed.
4.2 Excitation and Density
There are many different ways to infer density and several different definitions of
density. In addition to the trends in column density, the excitation temperature of
the gas can provide a measure of the density. It is defined by the relative column
densities of two energy levels of a given species,
−∆Eul N g kT ex u = u e ul . Nl gl
Here, u and l refer to the upper and lower levels, g is the statistical weight of the level, and ∆E is the energy difference between the levels. The populations of the levels are
133 Figure 4-1: This figure shows the relation between the column density of CO and H2. The current sample is shown as filled circles. Additional references appear in the figure legend.
determined from the balance between radiative and collisional processes. Stimulated
emission and absorption can usually be ignored when studying diffuse molecular gas
(Federman et al. 1984, 1994; Sonnentrucker et al. 2007; Goldsmith 2013).
4.2.1 H2 Excitation Temperatures
The relative populations of the J levels of molecular hydrogen can provide infor- mation about the physical environment of the gas. The photoexcitation of H2 begins with UV line absorption from the ground electronic state to an excited electronic
state. The excited H2 can then decay into a number of different vibrationally excited 134 levels of the electronic ground state, including the vibrational continuum that leads to photodissociation. At the low densities associated with diffuse gas, these vibra- tionally excited levels spontaneously decay to lower vibrational levels through electric quadruple transitions (Draine 2011). The radiative cascade from the high J levels is determined by the Einstein A coefficients. However, the low J levels have longer lifetimes, and collisional effects can play a larger role in the de-excitation. The low J levels of molecular hydrogen (J = 0 and 1) eventually come into thermal equilibrium with the surrounding gas and provide a measure of the kinetic temperature.
At lower kinetic temperatures, which is indicative of higher density gas associated with higher 12CO column densities, we consistently see detections of 13CO (Sheffer et al. 2007). This can be seen in Figure 4-2, which shows the H2 temperature his- togram for the current sample against the sample from Sheffer et al. (2007). The current sample is gray and can be found in Table 4.1. The Sheffer et al. (2007) sam- ple is in hashed red. Above the histogram are the temperature ranges from several additional data sets (Savage et al. 1977; Rachford et al. 2002; Burgh et al. 2007; Son- nentrucker et al. 2007; Sheffer et al. 2007). For each of the data sets, the ranges of temperatures with and without detections of 13CO are displayed above the histograms with 13CO detections in blue and non-detections in red. The symbols indicate the average value for each subset with dispersions (±1 σ) shown around each average.
In all the data sets we find 13CO detections at lower kinetic temperatures (∼20 K colder on average), indicative of higher density gas. Sheffer et al. (2007) attribute this connection between the kinetic temperature and detections of 13CO to the fact that because it is less abundant than 12CO, it is only observed deeper into the cloud where the densities are higher and the temperature lower. Section 4.5 discuses the implications of the isotopic CO ratio in more detail.
135 Table 4.1: H2 Excitation Temperatures
Sight Line Tex (K) Sight Line Tex (K) HD 108 61(11) HD 19820 61(8) HD 5689 64(8) HD 25443 56(11) HD 12882 63(11) HD 25638 63(12) HD 13841 60(7) HD 45314 67(19) HD 13969 58(8) HD 46223 53(10) HD 14053 61(7) HD 92964 63(8) BD+56◦ 0501 57(10) HD 97253 63(8) BD+56◦ 0563 66(8) HD 108927 60(7) HD 14443 48(4) HD 149404 48(5) HD 14476 55(9) HD 151932 60(10) BD+56◦ 0578 66(11) HD 168076 57(7) HD 14947 55(6) HD 170740 63(8) HD 15629 67(11) HD 192303 63(8) HD 16691 48(4) HD 192641 67(12) HD 17505 66(8) HD 193793 56(8) HD 17520 48(8) HD 197770 59(7) BD+60◦ 0586 47(4) HD 214419 56(30) HD 19243 66(9) HD 217086 54(6)
136 Figure 4-2: This figure shows the H2 temperature histogram for the current sample against the sample from Sheffer et al. (2007). The cur- rent sample is gray and the Sheffer et al. (2007) sample is in hashed red. Above the histogram are the temperature ranges from several additional data sets referenced in the figure legend. For each of the data sets, the ranges of temperatures with and without detections of 13CO are displayed above the histograms with 13CO detections in blue and non-detections of 13CO in red. The symbols indicate the average value for each subset with ±1 σ shown around each average. The H2 temperature comes from the ratio of the J = 0 and 1 levels.
137 4.2.2 CO Excitation Temperatures
13 In addition to H2 excitation temperatures and CO detections, the excitation analysis of CO provides a more direct density measure. Goldsmith (2013) used the excitation temperatures from CO with the kinetic temperature from H2 to model the density of collisional partners in diffuse molecular clouds. In particular, Sonnentrucker et al. (2007) and Goldsmith (2013) concluded that radiative excitation by nearby clouds does not play a major role in determining the excitation temperature of the lower rotational transitions of CO and that excitation is primarily by collisions. The excitation of CO and the inferred density is linked to the cloud structure and potential collision partners, electrons and atomic and molecular hydrogen. However, electrons have a much lower fractional abundance than atomic or molecular hydrogen at the densities and temperatures of diffuse molecular clouds, and thus should not be a significant source of collisional excitation for CO. Furthermore, the collisional cross sections of CO with H atoms are an order of magnitude smaller than those for H2 (Green & Thaddeus 1976; Shepler et al. 2007; Yang et al. 2013; Goldsmith et al.
2010a).
Goldsmith (2013) used the Meudon PDR code of Le Petit et al. (2006) to describe the transition from atomic to molecular hydrogen. He found that clouds with modest visual extinction (0.03 < AV < 1 mag) have a molecular hydrogen density that is two to ten time higher than the atomic hydrogen density, consistent with the transitional clouds studied by Kavars et al. (2005).
138 Combining the larger abundance and collisional cross section means that collisions
with H2 are likely the dominant source of excitation for CO molecules. Goldsmith (2013) adopted the rate coefficients and cross sections from Yang et al. (2013) to
model the excitation of CO by collisions with H2, assuming an ortho−to−para ratio
(OPR) equal to 3. Figure 4-3 shows the excitation of CO by collisions with H2 from his models for a larger set of kinetic temperatures, at our request.
Table 4.2 lists the CO temperatures from the Sheffer et al. (2008) sample. This
sample makes up the majority of the sight lines used by Goldsmith (2013) to cal-
culate the density from approximate kinetic temperatures. However, our analysis
recalculated the results using observed kinetic gas temperatures from H2 instead of a constant assumed kinetic temperature. We enlarge the sample by including our
results, but must first consider a different description of excitation temperature for
consistency. Table 3.13 lists the CO temperatures for the current sample and Table
4.3 shows the conversion of the temperatures to those needed for the analysis from
the following formulae, modified from Goldsmith (2013):
T10 = T01,
∆E21T02T01 T21 = , and ∆E20T01 − ∆E10T02
∆E32T03T02 T32 = . ∆E30T02 − ∆E20T03
These excitation temperatures are used with the corresponding kinetic temper- atures to find the density from each excited level with Figure 4-3. There are four kinetic temperatures shown in each panel of Figure 4-3, 20 K, 40 K, 60 K, 80 K. The y-axis is the CO excitation temperature between the labeled levels and the x-axis is the log of the density of H2 collisional partners. It is assumed that there are equal numbers of H and H2 so that the total density of collisional partners is a factor of
139 log10(2) larger. This figure has been adapted from Goldsmith (2013). Tables 4.4 and 4.5 show the values used to find the density inferred from each excitation temperature,
which is given along with the one sigma maximum and minimum values.
Figure 4-4 shows the density histograms from T CO(1-0) for the current sample and a combined comparable sample that includes all sight lines from the suite of results
in Goldsmith (2013) for the same range of N (CO) and N (H2) as in our observed sample. The current sample is the hashed red histogram and the Goldsmith (2013)
sample is the gray histogram. Above the histograms are the average value for each
data set with ±1 σ around each average. The figure indicates that the two samples
have a similar average density (within one sigma of log[nH+H2 ] ≈ 2.20).
Figure 4-5 shows the density histograms from T CO(2-1) for the current sample and a comparable sample from Goldsmith (2013). The samples are the same as in
Figure 4-4. The figure has very few data points from the current data set; however, it still indicates that the current sample also has an similar average density (within
one sigma of log[nH+H2 ] ≈ 2.75). There is only one detection of J = 3, toward HD
170740. The density from T CO(3-2) for HD 170740 is consistent with the other CO temperature measurements. However, there are too few other detections of J = 3 for densities from T CO(3-2) to be used to draw any conclusions about averages.
140 Table 4.2: Excitation Temperatures and Density from T CO(1-0) in Sheffer et al. (2008)
a b −3 Sight Line T H2 (1-0) T CO(1-0) log(n /cm ) (K) (K) BD+48◦ 3437 83 2.7 ≤ 0.80 BD+53◦ 2820 93 3.3 1.90 CPD −69 1743 79 2.7 ≤ 0.80 CPD −59 2603 77 3.0 1.60 HD 12323 82 3.1 1.75 HD 13268 92 3.4 1.90 HD 13745 66 4.0 2.30 HD 14434 99 4.4 2.35 HD 15137 104 3.1 1.60 HD 23478 55 3.4 2.10 HD 24190 66 3.1 1.80 HD 24398 57 3.4 2.05 HD 24534 54 4.6c 2.50 HD 27778 51 5.3 2.65 HD 30122 61 3.8 2.25 HD 36841 ... 2.7 ≤ 0.80c HD 37367 73 3.2 1.85 HD 37903 64 2.7 ≤ 0.80 HD 43818 ... 4.1 ≤ 2.65c HD 58510 90 2.9 1.30 HD 63005 78 3.6 2.10 HD 91983 61 2.7 ≤ 0.80 HD 93205 97 2.8 0.90 HD 93222 69 3.3 1.95 HD 93237 58 3.1 1.75 HD 93840 54 3.1 1.75 HD 94454 74 3.8 2.20 HD 96675 55 3.7 2.20 HD 99872 66 3.7 2.20 HD 102065 ... 3.6 ≤ 2.45c HD 106943 96 2.7 ≤ 0.80 HD 108002 77 3.2 1.80 HD 108639 88 3.0 1.50 HD 110434 87 2.7 ≤ 0.80 HD 112999 96 3.0 1.45 HD 114886 92 3.1 1.55
141 Table 4.2: —Continued
a b −3 Sight Line T H2 (1-0) T CO(1-0) log(n /cm ) (K) (K) HD 115071 71 3.7 2.15 HD 115455 81 2.9 1.40 HD 116852 66 3.2 1.85 HD 122879 90 2.9 1.35 HD 124314 74 3.8 2.20 HD 137595 72 3.9 2.25 HD 140037 ... 2.9 ≤ 1.50c HD 144965 70 4.3 2.30 HD 147683 58 5.2 2.55 HD 147888 44 13.6 3.10 HD 148937 69 3.7 2.15 HD 152590 64 4.1 2.35 HD 152723 76 4.0 2.30 HD 154368 47 3.0 1.65 HD 157857 86 4.6 2.40 HD 163758 79 4.0 2.30 HD 177989 49 3.3 2.15 HD 190918 102 2.7 ≤ 0.80 HD 192035 68 3.2 1.85 HD 195965 91 3.0 1.50 HD 198781 65 3.4 2.00 HD 203532 47 5.3 2.70 HD 208905 77 6.0 2.65 HD 209481 78 2.9 1.40 HD 209975 73 2.9 1.40 HD 210121 51 6.2d 2.75 HD 210809 87 3.1 1.70 HD 220057 65 3.0 1.60 HD 303308 91 3.1 1.70 HD 308813 73 3.8 2.25 a The T CO(1-0) excitation temperatures for HD 34078, HD 108610, and HD 200775 are not given in Sheffer et al. (2008) and are not included. b Refers to density of collision partners, (nH+nH2 ). c Corresponds to a T H2 (1-0) of 20 K, giving a maximum density for the given T CO(1-0) excitation temperature. d The T CO(1-0) excitation temperatures for HD 24534 and HD 210121 are from Sonnentrucker et al. (2007).
142 Table 4.3: CO Excitation Temperature Conversions
Sight Line T0−1 T0−2 T0−3 MIN T1−0 MAX MIN T2−1 MAX MIN T3−2 MAX (K) (K) (K) (K) (K) (K) (K) (K) (K) (K) (K) (K) HD 108 4.0 (0.1) 4.2 (0.3) ... 3.9 4.0 4.2 3.4 4.2 5.2 ... HD 5689 3.5 (0.3) ...... 3.2 3.5 3.9 ...... HD 12882 3.9 (0.4) ...... 3.5 3.9 4.3 ...... HD 13841 5.9 (1.2) ...... 4.7 5.9 7.0 ...... HD 13969 3.8 (0.8) ...... 3.0 3.8 4.6 ...... HD 14053 5.0 (2.0) ...... 3.1 5.0 7.0 ...... BD+56◦ 0501 4.0 (0.6) ...... 3.4 4.0 4.7 ...... BD+56◦ 0563 5.9 (2.3) ...... 3.6 5.9 8.3 ...... HD 14443 5.5 (2.0) ...... 3.5 5.5 7.5 ...... HD 14476 6.1 (2.8) ...... 3.4 6.1 8.9 ...... BD+56◦ 0578 6.4 (1.2) ...... 5.2 6.4 7.7 ...... HD 14947 3.0 (1.0) ...... 2.0 3.0 3.9 ...... HD 15629 5.0 (2.0) ...... 3.0 5.0 7.0 ...... HD 16691 3.1 (0.4) ...... 2.7 3.1 3.4 ...... HD 17505 3.3 (0.4) ...... 2.9 3.3 3.6 ...... HD 17520 5.7 (1.5) ...... 4.2 5.7 7.1 ...... BD+60◦ 0586 4.8 (0.6) 4.8 (0.3) ... 4.1 4.8 5.4 4.1 4.8 5.9 ... HD 19243 6.0 (1.0) ...... 5.0 6.0 6.9 ...... HD 19820 3.7 (0.7) ...... 2.9 3.7 4.4 ...... HD 25443 2.5 (0.2) ...... 2.3 2.5 2.7 ...... HD 25638 3.2 (0.6) ...... 2.6 3.2 3.9 ...... HD 45314 4.0 (0.3) 4.0 (0.1) ... 3.8 4.0 4.3 3.7 4.0 4.3 ... HD 46223 3.2 (0.2) 3.8 (0.1) ... 3.0 3.2 3.4 2.8 3.3 4.0 ... HD 92964 6.5 (1.5) ...... 4.9 6.5 8.0 ...... HD 97253 4.5 (0.2) ...... 4.3 4.5 4.7 ...... HD 108927 5.2 (0.3) 4.5 (0.2) ... 4.9 5.2 5.5 3.8 4.2 4.6 ... HD 149404 5.1 (0.5) 5.1 (0.2) ... 4.6 5.1 5.6 4.6 5.1 5.8 ... HD 151932 5.0 (0.8) 5.0 (0.5) ... 4.2 5.0 5.9 4.1 5.0 6.4 ... HD 168076 4.5 (0.7) ...... 3.8 4.5 5.2 ...... HD 170740 7.4 (0.6) 6.0 (0.3) 6.4 (0.3) 6.7 7.4 8.0 5.0 5.5 6.1 6.0 6.8 8.0 HD 192303 6.4 (2.3) ...... 4.1 6.4 8.8 ...... HD 192641 6.7 (3.8) ...... 2.9 6.7 10.5 ...... HD 193793 4.4 (0.5) 4.4 (0.5) ... 3.9 4.4 4.9 3.5 4.4 5.6 ... HD 197770 3.5 (0.5) 3.5 (0.2) ... 3.0 3.5 4.0 3.0 3.5 4.2 ... HD 214419 3.0 (0.5) ...... 2.5 3.0 3.5 ...... HD 217086 4.0 (0.9) ...... 3.1 4.0 5.0 ......
143 Table 4.4: Excitation Temperatures and Density from T CO(1-0)
a −3 Sight Line MIN T CO(1-0) MAX MIN T H2 (1-0) MAX MIN log(n /cm ) MAX HD 108 3.9 4.0 4.2 50 61 73 2.20 2.30 2.40 HD 5689 3.2 3.5 3.9 56 64 72 1.90 2.05 2.20 HD 12882 3.5 3.9 4.3 53 63 74 2.05 2.20 2.50 HD 13841 4.7 5.9 7.0 53 60 67 2.50 2.70 2.85 HD 13969 3.0 3.8 4.6 50 58 66 1.60 2.20 2.50 HD 14053 3.1 5.0 7.0 53 61 68 1.55 2.55 2.80 BD+56◦ 0501 3.4 4.0 4.7 47 57 68 2.00 2.30 2.40 BD+56◦ 0563 3.6 5.9 8.3 58 66 75 2.10 2.70 2.90 HD 14443 3.4 5.5 7.6 43 48 52 2.00 2.65 2.95 HD 14476 3.4 6.1 8.9 46 55 64 2.00 2.70 2.95 BD+56◦ 0578 5.2 6.4 7.7 55 66 78 2.55 2.70 2.85 HD 14947 2.0 3.0 3.9 49 55 60 ≤0.80 1.60 2.35 HD 15629 3.0 5.0 7.0 55 67 78 1.55 2.50 2.80 HD 16691 2.7 3.1 3.4 43 48 52 ≤0.80 1.85 2.15 HD 17505 2.9 3.3 3.6 58 66 75 1.30 1.90 2.20 HD 17520 4.2 5.7 7.1 40 48 55 2.45 2.70 2.85 BD+60◦ 0586 4.1 4.8 5.4 43 47 51 2.45 2.55 2.70 HD 19243 5.0 6.0 6.9 57 66 76 2.55 2.65 2.75 HD 19820 2.9 3.7 4.4 53 61 69 1.60 2.15 2.50 HD 25443 2.3 2.5 2.7 45 56 67 ≤0.80 ≤0.80 1.30 HD 25638 2.6 3.2 3.9 52 63 75 ≤0.80 1.90 2.20 HD 45314 3.8 4.0 4.3 47 67 86 2.10 2.30 2.45 HD 46223 3.0 3.2 3.4 43 53 63 1.50 1.90 2.10 HD 46223(13CO) 3.2 3.7 4.2 43 53 63 1.90 2.20 2.50 HD 92964 4.9 6.5 8.0 55 63 70 2.50 2.75 2.90 HD 97253 4.3 4.5 4.7 55 63 72 2.40 2.50 2.60 HD 108927 4.9 5.2 5.5 53 60 67 2.50 2.60 2.70 HD 108927(13CO) 4.9 5.2 5.5 53 60 67 2.50 2.60 2.70 HD 149404 4.6 5.1 5.6 44 48 53 2.50 2.65 2.70 HD 151932 4.2 5.0 5.9 50 60 70 2.35 2.55 2.75 HD 168076 3.8 4.5 5.2 50 57 63 2.15 2.45 2.60 HD 170740 6.7 7.4 8.0 55 63 70 2.70 2.85 2.90 HD 170740(13CO) 3.6 8.0 12.4 55 63 70 2.05 2.90 3.10 HD 192303 4.3 6.4 8.5 56 63 71 2.40 2.75 2.90 HD 192641 2.9 6.7 10.5 55 67 79 1.20 2.80 3.00 HD 193793 3.9 4.4 4.9 48 56 65 2.20 2.45 2.55 HD 197770 3.0 3.5 4.0 52 59 66 1.65 2.10 2.35 HD 214419 2.5 3.0 3.5 25 56 86 ≤0.80 1.60 2.40 HD 217086 3.1 4.0 5.0 49 54 60 1.85 2.30 2.55 a This refers to density of collisional parameters.
144 Table 4.5: Excitation Temperatures and Density from T CO(2-1) & T CO(3-2)
a −3 Sight Line MIN T CO(2-1) MAX MIN T H2 (1-0) MAX MIN log(n /cm ) MAX HD 108 3.4 4.2 5.2 50 61 73 1.75 2.50 2.90 BD+60◦ 0586 4.1 4.8 5.9 43 47 51 2.45 2.70 3.10 HD 45314 3.7 4.0 4.3 47 67 86 2.15 2.30 2.60 HD 46223 2.8 3.3 4.0 43 53 63 ≤ 0.80 1.90 2.50 HD 108927 3.8 4.2 4.6 53 60 67 2.15 2.50 2.75 HD 149404 4.6 5.1 5.8 44 48 53 2.70 2.85 3.00 HD 151932 4.1 5.0 6.4 50 60 70 2.40 2.80 3.10 HD 170740 5.0 5.5 6.1 55 63 70 2.70 2.90 3.10 HD 193793 3.5 4.4 5.6 48 56 65 2.00 2.60 3.00 HD 197770 3.0 3.5 4.2 52 59 66 1.40 2.05 2.55
a −3 Sight Line MIN T CO(3-2) MAX MIN T H2 (1-0) MAX MIN log(n /cm ) MAX HD 170740 6.0 6.8 8.0 55 63 70 2.10 2.60 3.15 a This refers to density of collisional parameters.
145 10
J=3-2 5
0 10
8
6 J=2-1 4
2
0
15
10
5 J=1-0
0 0 1 2 3
Figure 4-3: This figure shows models for the excitation temperatures (T ex) for the three lowest transitions of CO as a function of H2 density. The curves for each rotational transition indicate four different kinetic temperatures T k = 80 K, 60 K, 40 K, 20 K, from top to bottom. Figure adapted from Goldsmith (2013) with a larger set of kinetic temperatures.
146 Figure 4-4: This figure shows the density histograms from T CO(1-0) for the current sample and the combined sample from Goldsmith (2013). The current sample is the hashed red histogram and the Gold- smith (2013) sample is the gray histogram. Above the histograms are the average value for each data set with ±1 σ shown around each average.
147 Figure 4-5: This figure shows the density histograms from T CO(2-1) for the current sample and the literature sample from Goldsmith (2013). The current sample is the hashed red histogram and the Gold- smith (2013) sample is the gray histogram. Above the histograms are the average value for each data set with ±1 σ shown around each average.
148 4.2.3 C I Excitation Analysis
Assuming a steady state and detailed balance between the fine-structure levels, we can determine the approximate gas densities and temperatures of the gas containing neutral carbon from population ratios.
N n R (A + R + R ) + R R J=2 ≈ J=2 = 0,2 1,0 1,0 1,2 1,2 0,1 NJ=0 nJ=0 (A2,1 + R2,1 + R2,0)(A1,0 + R1,0 + R1,2) − R1,2 (A2,1 + R2,1) (4.1) and
N n R (A + R + R ) + R (A + R ) J=1 ≈ J=1 = 0,1 2,1 2,1 2,0 0,2 2,1 2,1 . NJ=0 nJ=0 (A2,1 + R2,1 + R2,0)(A1,0 + R1,0 + R1,2) − R1,2 (A2,1 + R2,1) (4.2)
Au,l are the spontaneous decay rates between the upper level u and the lower level l.
P k k Ri,j are k γi,jnk + Gi,j, where γi,j represent the collisional rate coefficients involving species k and Gi,j are the photon pumping rates between fine-structure levels and excited electronic levels. However, photon pumping was not included in the anal- ysis because it had a negligible effect on the populations when the radiation field is comparable to the interstellar average. If the neutral carbon containing gas has approximately uniform density and composition, we can approximate the gas density ratio to be the column density ratio, as described in the work of Zsarg´o& Federman
(2003).
Figure 4-6 shows the density of collisional partners derived from the C I ratios toward HD 108. The results for the total line of sight column density are shown in the
first panel with the star name. The remainder of the panels indicate the results for each velocity component. Within each panel the solid lines correspond to the 3 sigma maximum and minimum density derived for the J =1/J =0 ratio and the dotted lines correspond to the 3 sigma maximum and minimum density derived for the J =2/J =0 ratio. If only a single line of the pair is shown, it indicates an upper limit for that
149 measurement. The horizontal lines in the figure show the one sigma uncertainty in
the kinetic temperature from H2 observations. Similar figures for other sight lines are shown in Appendix F.
In Zsarg´o& Federman (2003), the density ranges from the different ratios overlap;
however, the column densities for the current sample are much higher and the cor-
responding errors are on the order of 1% of the column density instead of 10% as in
Zsarg´o& Federman (2003). As such, the derived densities from the N (J =1)/N (J =0)
and N (J =2)/N (J =0) ratios often give different density ranges.
The range in density for each velocity component is determined from the
N (J =1)/N (J =0) ratio at the largest value of the one sigma kinetic temperature
range and the N (J =2)/N (J =0) ratio at the smallest value. The range in density for the line of sight is determined from the minimum and maximum density from each component, weighted by the N (J = 0) for that component. Table 4.6 shows the
inferred density ranges for the directions with HST spectra.
Overall we see that the densities derived from the N (J =2)/N (J =0) ratio tend
to be larger than those derived from the N (J =1)/N (J =0) ratio. Although this is
not seen in Zsarg´o& Federman (2003), this is expected as the material containing
observable amounts of J = 2 should be deeper into the cloud where there tends to be
higher densities. A detailed comparison of inferred densities for the current sample
is presented in the next chapter.
150 Table 4.6: C I Densities
Star Vlsr nH+H2 min max (km s−1) (cm−3) (cm−3) HD 108 -16.8 10 85 -9.7 5 45 -5.5 35 75 -1.7 5 82 2.4 5 82 6.6 35 150 Weighteda Total 24 90 Combinedb Total 28 70
HD 13841 -41.2 5 210 -15.6 ... 320 -10.5 15 180 -2.9 5 55 Weighteda Total 9 150 Combinedb Total 10 90
HD 25443 -10.9 16 49 -7.2 5 25 1.3 7 33 6.4 9 32 Weighteda Total 10 36 Combinedb Total 11 30
HD 46223 -5.8 ... 11 1.4 6 32 6.6 20 115 8.7 2 9 14.4 19 65 Weighteda Total 6 26 Combinedb Total 7 21
151 Table 4.6: —Continued
Star Vlsr nH+H2 min max (km s−1) (cm−3) (cm−3) HD 108927c 18.2 20 42 Weighteda Total 20 42 Combinedb Total 21 43
HD 170740 -12.5 ... 57 -9.2 15 105 Weighteda Total 15 94 Combinedb Total 15 87 aThis total is derived from the density range of each component weighted by its fraction of the total J = 0 column density. bThis total comes from using total column densities from the J levels to determine the ratios and density ranges. cHD 108927 has a second minor component but the resulting allowed density ranges do not constrain the results.
152 Figure 4-6: This figure shows C I density of collisional partners versus tem- perature toward HD 108.
153 4.3 Chemistry and Density
The chemical analysis is based largely on previous models involving CN, CH, C2, and NH. We primarily follow the work of Federman et al. (1994), who used steady- state chemistry (where the production and destruction rates are balanced) leading to CN to estimate the local gas density. The gas phase chemical network for CN production that is used here has been described in a series of papers (Federman et al.
1984; Federman & Lambert 1988; Federman & Huntress 1989; Federman et al. 1994;
Pan et al. 2001).
4.3.1 CN Chemistry
The most important pathways for CN production are shown in Figure 4-7. The local abundance of CN is dependent on the amounts of CH, C2, and NH found at the various depths throughout the diffuse cloud. The optical depth is a measure of the distance into a cloud in units of the mean free path and can vary with wavelength.
τUV is the optical depth of interstellar grains at 1000 A,˚ chosen to be in the middle of the range for photodissociating C2 and CN − Lyman limit to 1200 A˚ (Federman et al. 1994). The calculations for values of τUV follow the work of Federman et al. (1994) and depend on the grain properties,
τUV = (2)RV E(B −V),
where RV is the ratio of total to selective extinction (A(V )/E(B−V )) and is typically
between 4 and 6 in dark clouds but is smaller in the diffuse ISM with RV = 3.1(0.2) (Whittet et al. 1976), which is the value being used in the calculations. The typical
value of the prefactor in the τUV equation (two) is based on extinction curves of Code et al. (1976), corrected for the forward scattering of dust at far UV wavelengths
154 (Federman et al. 1994). The majority of our sight lines have typical grain properties
and use the initial value of the prefactor. However, sight lines with enhanced UV
extinction have larger values, while sight lines with low UV extinction have smaller
values. The final value of the prefactor was modified to account for the differences
in the total to selective extinction and ultraviolet extinction. Uncertainties in τUV lead to uncertainties up to ≈ 30% in the inferred density (Sheffer et al. 2008). Our values of τUV are given in Table 4.7 along with the values of the prefactor used in the conversion and the E(B-V ) averages of the values given in the corresponding references.
The dominant production and destruction pathways change with optical depth into a cloud. At low optical depths (up to τUV ∼ 3), the production of CN is dominated by the ion-molecule reaction between C+ and NH. C+ combines with NH through
three channels:
C+ + NH → CN+ + H,
C+ + NH → CN + H+, and
C+ + NH → CH+ + N.
However, only the first channel (CN+ + H) is fast enough to contribute significantly
(Viggiano et al. 1980; Federer et al. 1986). Reactions with CH+ can also produce
CN+ but are secondary,
CH+ + N → CN+ + H.
155 Table 4.7: τUV Determinations
Sight Line E(B-V ) Prefactor τUV References HD 108 0.49 2.0 3.04 1,2,3,4 HD 5689 0.66 1.7 3.48 1,2,3 HD 12882 0.38 2.0 2.36 1 HD 13841 0.43 2.0 2.67 1,2,5 HD 13969 0.56 2.0 3.47 2,3,4,5,6 HD 14053 0.49 2.0 3.04 1,2,5,7 BD+56◦ 0501 0.52 2.0 3.22 8,7 BD+56◦ 0563 0.58 2.5 4.50 2,5 HD 14443 0.53 2.5 4.11 1,2,5 HD 14476 0.63 3.0 5.86 2,5 BD+56◦ 0578 0.55 2.0 3.41 2,5 HD 14947 0.77 2.0 4.77 1,2,3,4,5,6,9,10,11 HD 15629 0.73 2.0 4.53 1,2,5,12,13 HD 16691 0.80 2.0 4.96 2,3,10,14 HD 17505 0.70 2.5 5.43 1,2,5,7,13,15 HD 17520 0.61 2.0 3.78 1,2,5 BD+60◦ 0586 0.61 1.9 3.50 1,2,3,16,17 HD 19243 0.47 2.0 2.91 1 HD 19820 0.78 2.0 4.84 3,4,5,7 HD 25443 0.60 2.5 4.65 3,7,18 HD 25638 0.69 2.0 4.28 7 HD 45314 0.44 1.7 2.32 3,4,5,14,15,19 HD 46223 0.53 2.0 3.29 1,5,10 HD 92964 0.39 2.0 2.42 1,5,20 HD 97253 0.48 2.0 2.98 3,4,5,21 HD 108927 0.23 2.0 1.43 22 HD 149404 0.66 2.0 1.89 3,7,15 HD 151932 0.43 2.0 2.66 23 HD 168076 0.76 2.0 4.36 3,4,5,7,10,15 HD 170740 0.48 2.0 2.98 24 HD 192303 0.54 2.0 3.35 1 HD 192641 0.65 2.0 4.03 23 HD 193793 0.88 2.0 5.46 23 HD 197770 0.52 2.5 4.03 3,7,15,25 HD 214419 0.51 2.0 3.16 23 HD 217086 0.94 2.0 5.83 2,3,4,5,7,10,14,26
156 Table 4.7: —Continued
1Aiello et al. (1988) — A(V ) 2Buss et al. (1994) — E(B-V ) 3Gordon et al. (2009) — E(B-V ), E(15-V ) 4Savage et al. (1985) — E(B-V ), E(15-V ) 5van Breda & Whittet (1977) — E(B-V ) 6Wegner (2002) — E(B-V ) 7Whittet et al. (1976) — A(V ), E(B-V ) 8Hanson & Clayton (1993) — E(B-V ) 9Garmany & Stencel (1992) — E(B-V ), E(15-V ) 10Humphreys (1978) — E(B-V ) 11Patriarchi et al. (2003) — E(B-V ) 12Franco et al. (1985) — A(V ) 13Papaj et al. (1991) — E(15-V ) 14The et al. (1989) — A(V ), E(B-V ) 15Zdanaviˇcius& Zdanaviˇcius(2002) — E(B-V ) 16Massa & Savage (1984) — E(B-V ) 17Papaj & Krelowski (1992) — A(V ), E(B-V ) 18Patriarchi & Perinotto (1999) — A(V ) 19Krelowski & Strobel (1983) — E(B-V ) 20Fitzpatrick & Massa (2007) — A(V ) 21Jenniskens & Greenberg (1993) — E(B-V ) 22Andersson et al. (2002) — E(B-V ) 23Vacca & Torres-Dodgen (1990) — A(V ), E(B-V ) 24Rachford et al. (2009) — E(B-V ) 25Wegner (2003) — E(B-V ) 26Geminale & Popowski (2004) — E(B-V ), E(15-V )
157 Table 4.8: CN Chemistry of Sight Lines in the Current Sample
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 108 -5.1 4.1 ... 4.4 1.5 1.5 400 6.9 1.4 ... 2.9 1.1 1.1 975 Total 650 HD 5689 -12.6 9.3 ... 16.0 ≤ 4.4 ≤ 4.4 ≤ 450 HD 12882 0.2 7.6 ... 19.2 ≤ 2.7 ≤ 2.7 ≤ 425 158 HD 13841 -41.6 4.8 ... 4.9 ≤ 0.7 ≤ 0.7 ≤ 175 HD 13969 -10.8 4.7 ... 4.2 ≤ 1.6 ≤ 1.6 ≤ 300 HD 14053 -8.6 9.4 ... 7.8 ≤ 2.3 ≤ 2.3 ≤ 275 BD+56◦ 0501 -9.2 5.0 ... 6.6 ≤ 2.8 ≤ 2.8 ≤ 600 BD+56◦ 0563 -8.4 36.0 ... 3.0 ≤ 1.5 ≤ 1.5 ≤ 225 HD 14443 -6.2 3.4 ... 4.0 ≤ 1.2 ≤ 1.2 ≤ 225 HD 14476 -0.6 3.7 ... 4.2 ≤ 2.2 ≤ 2.2 ≤ 200 BD+56◦ 0578 -12.4 5.1 ... 9.5 ≤ 3.2 ≤ 3.2 ≤ 675 HD 14947 -7.7 11.0 ... 5.4 ≤ 1.7 ≤ 1.7 ≤ 75 HD 15629 -5.1 5.5 ... 5.7 ≤ 3.6 ≤ 3.6 ≤ 450 HD 16691 -12.7 9.5 ... 10.5 ≤ 8.3 ≤ 8.3 ≤ 1075 Table 4.8: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 17505 -4.3 7.8 ... 3.9 ≤ 1.4 ≤ 1.4 ≤ 55 HD 17520 -2.5 4.7 ... 4.8 ≤ 2.2 ≤ 2.2 ≤ 375 BD+60◦ 0586 -2.2 6.4 ... 0.0 ≤ 1.7 ≤ 1.7 ≤ 225 HD 19243 0.5 4.9 ... 3.0 ≤ 0.8 ≤ 0.8 ≤ 175 HD 19820 -14.9 17.2 ... 8.2 ≤ 2.5 ≤ 2.5 ≤ 65 HD 25443 -8.1 4.6 ... 3.3 ≤ 1.6 ≤ 1.6 ≤ 150
159 HD 25638 2.4 10.5 ... 8.5 3.8 3.9 225 5.7 11.1 ... 5.8 2.0 2.1 100 Total 175 HD 45314 1.7 6.1 ... 7.4 ≤ 1.8 ≤ 1.8 ≤ 350 HD 46223 11.3 9.7 ... 4.9 ≤ 1.3 ≤ 1.3 ≤ 125 HD 92964 -9.8 2.2 ... 0.4 0.05 0.04 50 -6.2 6.3 ... 1.8 0.3 0.2 75 -0.2 4.1 ... 2.1 0.4 0.4 125 Total 100 HD 97253 -14.6 13.4 21.4 10.6 3.3 3.8 250 -14.6 13.4 21.4 12.1 3.3 3.9 225 -7.3 4.2 ... 1.2 0.2 0.1 75 Total 225 Table 4.8: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 108927 1.5 14.1 13.8 11.8 2.0 2.2 250 HD 149404 -11.8 5.9 ... 4.0 0.7 0.7 125 7.4 14.2 ≤ 20.0 18.1 3.6 3.5 275 Total 250 HD 151932 0.1 4.7 ... 0.7 0.1 0.1 50
160 6.6 26.0 30.2 15.1 2.6 3.2 175 Total 175 HD 168076 0.9 9.1 21.1 8.3 1.8 2.1 225 6.5 2.4 4.0 1.7 0.3 0.4 175 Total 225 HD 170740 6.6 16.0 25.7 18.7 7.8 9.1 450 HD 192303 8.9 5.6 ... 8.0 ≤ 3.8 ≤ 3.8 ≤ 750 HD 192641 4.0 9.6 ... 4.7 ≤ 1.5 ≤ 1.5 ≤ 100 HD 193793 6.2 7.3 ... 2.4 ≤ 0.7 ≤ 0.7 ≤ 25 HD 197770 -1.4 31.5 26.9 20.0 7.0 8.3 150 HD 214419 -19.5 8.6 ... 12.7 5.6 5.6 750 HD 217086 -15.7 16.1 ... 6.9 ≤ 2.5 ≤ 2.6 ≤ 25 Table 4.9: CN Chemistry of Sight Lines from the Sheffer et al. (2008) Sample
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) BD+48◦ 3437 5.2 4.8 ... 3.1 0.6 0.6 175 BD+53◦ 2820 0.9 1.6 ... ≤ 3.3 ≤ 0.9 ≤ 0.9 ≤ 725 6.9 2.5 ... ≤ 3.8 ≤ 0.9 ≤ 0.9 ≤ 475 Total ≤ 600 HD 12323 -13.5 1.8 ... 3.2 0.9 0.9 625 -9.7 3.4 ... 4.3 1.0 1.0 400 Total 500 161 HD 13268 -10.4 2.7 ... 2.5 0.6 0.6 325 -16.4 1.2 ... ≤ 1.5 ≤ 0.4 ≤ 0.4 ≤ 475 -1.0 1.8 ... ≤ 1.7 ≤ 0.4 ≤ 0.4 ≤ 325 -7.4 5.4 ... 5.3 1.3 1.4 350 Total 350 HD 13745 -18.1 4.0 ... ≤ 3.2 ≤ 0.6 ≤ 0.6 ≤ 325 -43.9 5.9 ... 4.2 0.7 0.8 275 Total 275 HD 14434 -1.0 9.3 ... 2.9 0.6 0.6 70 HD 15137 -0.2 2.6 ... ≤ 1.5 ≤ 0.4 ≤ 0.4 ≤ 200 -13.4 2.2 ... ≤ 1.5 ≤ 0.4 ≤ 0.4 ≤ 225 -7.4 1.6 ... ≤ 1.3 ≤ 0.4 ≤ 0.4 ≤ 325 Total ≤ 250 HD 22951 7.1 12.0 3.6 3.9 0.6 0.6 80 Table 4.9: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 23180 4.6 7.0 ≤ 4.0 ≤ 2.9 ≤ 0.3 ≤ 0.3 ≤ 70 7.3 12.0 23.0 9.8 1.7 1.9 150 Total 150 HD 23478 4.1 13.5 7.8 6.5 1.0 1.1 175 7.7 4.7 6.2 4.7 0.8 0.9 250 Total 175
162 HD 24190 6.5 9.5 ... 4.7 0.8 0.8 100 HD 24398 6.8 22.0 35.0 34.3 3.9 4.1 250 HD 24534 7.1 25.9 31.0 22.6 6.6 7.9 150 5.0 6.0 3.4 3.4 0.7 0.6 75 Total 125 HD 27778 5.0 12.8 24.0 19.8 8.9 10.2 225 7.2 9.4 14.0 16.1 5.1 4.1 275 Total 250 HD 30122 6.9 15.7 ... 7.6 1.6 1.6 150 HD 36841 10.6 9.9 ... 5.5 1.6 1.6 175 HD 37367 3.8 3.2 ... ≤ 1.6 ≤ 0.3 ≤ 0.3 ≤ 150 6.2 9.8 ... ≤ 2.5 ≤ 0.3 ≤ 0.3 ≤ 75 Total ≤ 125 Table 4.9: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 43818 1.2 2.8 ... ≤ 1.8 ≤ 0.6 ≤ 0.6 ≤ 225 -3.9 2.7 ... ≤ 1.8 ≤ 0.6 ≤ 0.6 ≤ 225 5.2 2.0 ... ≤ 1.7 ≤ 0.6 ≤ 0.6 ≤ 325 -7.0 3.8 ... ≤ 2.0 ≤ 0.6 ≤ 0.6 ≤ 175 Total ≤ 225 HD 58510 29.6 5.1 ... ≤ 2.8 ≤ 0.5 ≤ 0.5 ≤ 125 HD 63005 14.3 4.6 ... 6.7 1.4 1.4 425 163 21.0 4.4 ... ≤ 5.2 ≤ 1.0 ≤ 1.0 ≤ 325 Total 425 HD 96675 4.1 22.8 ... 31.6 6.3 6.2 400 HD 99872 3.2 12.6 ... ≤ 4.3 ≤ 0.6 ≤ 0.4 ≤ 75 HD 102065 1.0 1.1 ... ≤ 1.7 ≤ 0.6 ≤ 0.4 ≤ 1075 3.8 6.0 ... ≤ 5.1 ≤ 0.9 ≤ 1.0 ≤ 525 Total ≤ 750 HD 115455 -3.3 17.0 ... ≤ 6.6 ≤ 1.6 ≤ 1.7 ≤ 125 HD 137595 5.4 12.2 ... ≤ 5.0 ≤ 0.8 ≤ 0.8 ≤ 125 HD 147888 Total 21.8 16.6 18.2 2.1 2.0 150 HD 147993 Total 14.9 26.0 17.4 2.1 2.4 225 HD 148184 Total 34.0 35.0 9.6 1.3 1.4 75 HD 149757 Total 25.0 18.0 20.6 2.6 2.2 175 Table 4.9: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 154368 -13.1 2.1 ... 0.8 0.2 0.2 50 3.3 54.1 51.0 46.0 27.0 28.5 225 Total 225 HD 157857 0.0 5.0 ... ≤ 4.3 ≤ 1.2 ≤ 1.2 ≤ 325 4.2 2.8 ... ≤ 3.6 ≤ 1.2 ≤ 1.2 ≤ 550 164 Total ≤ 425 HD 185418 6.8 13.0 ≤ 10.0 ≤ 18.1 ≤ 0.5 ≤ 0.6 ≤ 25 HD 190918 18.3 1.2 ... ≤ 1.2 ≤ 0.3 ≤ 0.3 ≤ 400 2.1 1.7 ... ≤ 1.4 ≤ 0.3 ≤ 0.3 ≤ 275 Total ≤ 350 HD 192035 1.4 3.3 ≤ 5.1 ≤ 2.9 ≤ 0.5 ≤ 0.6 ≤ 200 5.6 11.4 17.7 20.0 5.1 4.5 450 9.4 2.6 ≤ 4.1 ≤ 2.8 ≤ 0.5 ≤ 0.6 ≤ 250 Total 0 HD 192639 7.3 28.0 ≤ 10.0 ≤ 39.4 ≤ 0.7 ≤ 0.8 ≤ 25 HD 198781 5.4 13.2 ... 10.1 3.3 3.4 275 Table 4.9: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 206267 -0.9 2.1 7.9 1.6 0.9 1.0 225 0.9 9.0 4.4 2.6 0.5 0.4 75 -2.7 8.8 61.1 7.6 7.0 7.7 275 -4.7 4.1 2.6 1.3 0.3 0.4 75 165 Total 250 HD 207198 -0.2 9.8 5.9 3.7 0.9 1.1 75 -2.1 10.1 15.6 5.9 2.4 2.7 150 -6.2 5.3 8.5 3.3 1.3 1.6 175 Total 150 HD 208440 Total 11.7 ... ≤ 5.6 ≤ 0.9 ≤ 1.0 ≤ 125 HD 208905 Total 5.4 ... ≤ 4.9 ≤ 0.9 ≤ 0.9 ≤ 275 HD 209399 Total 7.9 ... ≤ 5.1 ≤ 0.9 ≤ 0.9 ≤ 275 HD 209481 -1.1 4.2 ... 0.7 0.5 0.1 50 HD 209975 Total 8.5 ... ≤ 5.2 ≤ 0.9 ≤ 0.9 ≤ 200 Table 4.9: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 210121 Total 28.6 65.0 316.0 17.4 19.9 225 HD 210809 -0.6 1.6 ... ≤ 2.9 ≤ 0.6 ≤ 0.6 ≤ 425 166 3.4 3.9 ... ≤ 3.3 ≤ 0.6 ≤ 0.6 ≤ 175 Total ≤ 300 HD 210839 -0.4 7.9 4.4 4.6 0.9 0.9 150 -2.3 12.9 12.6 12.3 2.6 2.8 275 Total 225 HD 220057 -1.8 9.3 ... 7.7 1.4 1.4 250 1.9 3.9 ... ≤ 3.6 ≤ 0.7 ≤ 0.7 ≤ 275 Total 250 Table 4.10: CN Chemistry of Additional Sight Lines from the Pan et al. (2004) Sample
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 37903 Total 9.2 ≤ 9.9 4.5 0.8 0.8 100 HD 203374 -1.2 6.9 ... 3.4 1.0 1.0 125 1.2 4.5 ... 1.6 0.4 0.3 75 167 3.1 6.3 ... 2.8 0.8 0.8 100 HD 204827 -9.2 7.0 ... 3.0 1.3 1.3 20 -6.4 5.3 ... 2.2 0.9 1.0 20 -4.3 31.4 ... 25.1 17.1 15.4 75 -2.0 13.6 ... 10.2 5.8 5.8 55 0.5 26.9 ... 20.6 12.3 12.1 60 HD 206165 -1.3 5.0 3.5 3.1 0.7 0.8 175 0.3 5.7 6.0 4.8 1.2 1.3 275 3.7 4.1 2.5 2.4 0.5 0.6 175 HD 206183 -1.6 7.8 ... 4.6 1.0 1.0 175 Table 4.10: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 206773 -0.9 4.5 ... 2.2 0.5 0.5 125 HD 207260 -8.1 5.1 ... 2.5 0.5 0.4 100 HD 207308 -2.4 18.4 ... 24.1 9.0 8.9 550 -0.5 8.6 ... 7.1 2.0 2.1 275 168 HD 207538 -4.4 7.0 ... 2.6 0.6 0.5 75 -2.5 11.7 ... 4.8 1.5 1.4 85 -0.5 11.8 ... 5.6 1.8 1.7 100 1.9 4.4 ... 2.5 0.9 0.9 125 HD 208266 -4.4 13.0 ... 13.2 4.4 4.5 375 -2.8 12.4 ... 4.0 0.9 0.9 85 -0.6 4.0 ... 1.7 0.4 0.4 125 HD 208501 -2.0 17.4 ... 14.7 7.5 7.5 225 -0.3 9.1 ... 4.5 1.7 1.7 85 1.5 19.4 ... 11.7 4.5 4.8 125 Table 4.10: —Continued
Sight Line V lsr N o(CH) N o(C2) N p(C2) N o(CN) N p(CN) nH+H2 (km s−1) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (1012cm−2) (cm−3) HD 210839 -2.3 12.9 12.6 12.3 2.6 2.8 275 169 -0.4 7.9 4.4 4.6 0.9 0.9 150 HD 216532 3.4 11.2 ... 11.7 3.1 3.1 375 HD 216898 4.8 5.4 ... 2.5 0.8 0.8 125 HD 217035 A Total 16.8 ... ≤ 4.8 ≤ 0.9 ≤ 0.9 ≤ 85 HD 217312 0.0 1.6 ... 1.4 0.4 0.4 325 5.2 4.2 ... 2.1 0.5 0.5 175 7.2 8.1 ... 2.6 0.6 0.6 100 Figure 4-7: This figure shows the most important pathways for the produc- tion of CN in diffuse molecular clouds. Original figure from Fe- derman et al. (1994).
Once CN+ is produced, it reacts with molecular hydrogen (hydrogen abstraction)
+ + to form HCN and H2CN ,
+ + CN + H2 → HCN + H and
+ + HCN + H2 → H2CN + H.
Each of these molecules can undergo electron dissociative recombination to form CN:
HCN+ + e → CN + H and
+ H2CN + e → CN + H2.
At larger optical depths (3 ≤ τUV ≤ 4.5), neutral-neutral reactions begin to be
170 more important and become the main the CN production pathways,
C2 + N → CN + C and CH + N → CN + H.
The neutral-neutral reactions vary weakly with temperature and typically have a temperature dependence in their rate constants of (T /300)0.5 based on velocity, while ion-neutral reactions (where the neutral has a large dipole moment such as CH and
NH) vary inversely with temperature (Marquette et al. 1985; Rebrion et al. 1988).
As the optical depth increases further (τUV ≥ 4.5), the column density of CN begins to plateau as photochemical destruction now competes with gas phase chemical destruction. These chemical production and destruction pathways lead to the steady state rate equations of Federman et al. (1994),
+ N(CN) = n[k3x(N)N(C2)+k4x(N)N(CH)+k5x(C )N(NH)α] , and G(CN)+k6+x(O)n + nα[k1x(C )N(CH)] N(C2) = . G(C2)+k2+x(O)n+k3x(N)n
The density here refers to the total proton gas density (nH+2nH2 ). In diffuse gas, the
proton gas density is 1.5 times the density of collisional partners (nH+nH2 ), where we have assumed the number density of atomic and molecular hydrogen are approxi- mately equal, as previously discussed in Section 4.2.2. In the preceding equations, the kx are reaction rate constants (see Federman et al. (1994)) and alpha (α) is a factor that accounts for the conversion of C+ into CO (see Federman & Huntress (1989)).
The photodestruction rate (G) is given by
−τUV G = IUVG0e .
G0 and IUV are the photodissociation rate at the cloud surface and the strength of the interstellar UV radiation field.
171 Knowing the column densities of CH, C2, and NH (when available) allows us to use the steady state analytical expressions from Federman et al. (1994) to derive the gas density of the CN rich gas. Often, NH is not observed. Fortunately, there is a roughly linear relationship between NH and CH from which we can use the observed N (CH) to determine a value for N (NH). The CH/NH ratio is 0.085 based on NH surveys; see Welty et al. (2006) for a compilation. The values for other species are based on average abundances for similar diffuse clouds. The fractional abundances of ionized carbon, atomic nitrogen, and atomic oxygen are 1.4x10−4, 7.5x10−5, and 3.2x10−4, respectively (Federman et al. 1994; Pan 2002; Sheffer et al. 2008). In the majority of cases, the average interstellar radiation field is used to model the observed molecular column densities. The rate coefficients come from the works of Knauth (2001), Pan et al. (2001), and McElroy et al. (2013).
Tables 4.8, 4.9, and 4.10 show the observed and predicted column densities from the CN chemistry (designated by an o or p subscript) for CH, C2, and CN as well as the gas density of collisional partners for the current sample, the Sheffer et al. (2008) sample, and the Pan et al. (2004) sample. Figure 4-8 shows the density histograms from CN chemistry for detections and upper limits in the current sample and the combined sample from Pan et al. (2004) and Sheffer et al. (2008). While there are very few three sigma detections of CN in the current sample, the average density of the two samples are indistinguishable when the upper limits from the CN non-detections are included.
4.3.2 CH+ like CH Chemistry
For the majority of sight lines in our sample, the total CH is found to resemble
CH+ like gas instead of CN like gas. When a significant amount of CH+ is present, the production and destruction routes of CH from CH+ can be combined with the observed column density of CH+ and CH present to determine the gas density (Zsarg´o
172 Figure 4-8: This figure shows the density histograms from CN chemistry for detections and upper limits in the current sample and the com- bined literature sample from Pan et al. (2004) and Sheffer et al. (2008). The current sample is the hashed histogram and the lit- erature sample is the gray histogram. Above the histograms are
the average value (log[nH+H2 ] ≈ 2.25) for each data set with ±1 σ shown around each average.
& Federman 2003). Assuming steady state, the most important reactions are
+ + CH + H2 → CH2 + H,
+ + CH2 + H2 → CH3 + H, and
+ CH3 + e → CH + H2 or 2H.
The rate equation for CH under these circumstances from Welty et al. (2006) can be
173 written as 0.67 k(CH+, H ) N(CH+) f(H ) n N(CH) = 2 2 , (4.3) 2 IUV G(CH)
+ The factor 0.67 represents the fraction of dissociative recombinations of CH3 pro- ducing CH (Herbst 1978; Vejby-Christensen et al. 1997). The molecular hydrogen fraction, f(H2), is defined as 2N(H2)/[N(H I)+2N(H2)]. The rate coefficient for the initial reaction k is 1.2 × 10−9 cm3 s−1 (Viggiano et al. 1980; Federer et al. 1986).
The CH+ like CH is from less dense gas than CN like CH gas and the lower density chemistry involving CH+ is less efficient at producing 12CO, as noted above in Section
4.1. This is discussed further in Section 4.4.
4.4 h and χ Persei
The sight lines in h and χ Per form a significant subgroup for the current sample.
They are useful for not only looking at small scale variations but investigating the local physical conditions in more detail. Understanding the physical conditions is critical for these sight lines in particular, as they show a deviation in some of the established trends between species seen with the other sight lines in the current sample.
4.4.1 CO versus H2
Figure 4-1 showed a plot of N (CO) versus N (H2) demonstrating a change in the CO photochemistry as CO starts to self shield. Figure 4-9 highlights the sight lines in h and χ Per by showing them as filled points. There are several points in h and
χ Per that come from Sheffer et al. (2008); however, our analysis is based on all the available sight lines in h and χ Per. The h and χ Per points are evenly distributed about the trend line.
The column density of CH is linearly correlated with H2 as shown in Figure 4-10 with the earlier results from Pan et al. (2005) and Sheffer et al. (2007) indicated. 174 Figure 4-9: This figure shows the relation between the column density of CO and H2. The current sample is indicated by the circles (both filled and unfilled). Filled symbols indicate the sight line is from the h and χ Per subsample.
Previous results are evenly distributed about the linear trend, shown as a solid line.
The current sample is indicated by the circles (both filled and unfilled). The filled symbols (both circles and triangles) indicate the sight line is member of the h and χ
Per subsample from all studies and show a marked deviation from the linear trend.
Our remaining sight lines, the unfilled circles, also lie about the average line. In particular, there seems to be less CH compared to the amount of H2 present. This is also seen in other figures.
The column densities of CO and CH in Figure 4-11 show a similar trend and break
175 to that of N (CO) versus N (H2) in Figure 4-9. Unlike the h and χ Per subsample in Figure 4-9, the h and χ Per subsample differs from the results of the other sight lines. There is less CH for the amount of CO present in the clouds toward h and χ
Per. The departure can be explained by either an increased local radiation field or a decreased local gas density. We can now explore the cause in more detail.
Figure 4-10: This figure shows the linear relation between the column density of CH and H2. The current sample is indicated by the circles (both filled and unfilled). Filled symbols indicate the sight line is from the h and χ Per subsample.
176 Figure 4-11: This figure shows the relation between the column density of CO and CH. The current sample is indicated by the circles (both filled and unfilled). Filled symbols indicate the sight line is from the h and χ Per subsample.
177 4.4.2 The CN/CH+ and CH/CH+ Ratios
There are several different ways to probe the density of the gas. The column density ratio of two species existing in material with different densities can provide a qualitative measure of the gas density. Typically, we use a species present in denser gas like CN or CH (when CN is not detected) and a species present in more diffuse gas such as CH+.
Figure 4-12: This figure shows the N (CN)/N (CH+) ratio on the N (CO) ver- sus N (H2) plot for the current sample and a combined literature sample from Pan et al. (2005) and Sheffer et al. (2008). The current sample is in red and the comparison sample is in black. Above the plot is a key relating the size of the symbols to the values of the ratio.
178 + Figure 4-12 shows the N (CN)/N (CH ) ratio on the N (CO) versus N (H2) plot for the current sample and previous results. The current sample is in red and earlier
results are in black. Above the plot is a key relating the size of the symbols to
the values of the ratio. The majority of the current sample tends to have smaller
N (CN)/N (CH+) ratios. Thus on average our sample probes lower density gas even
though our data are distributed equally about the average in the N (CO) versus N (H2) plot, including the h and χ Per subsample. When CN was not detected an upper limit was used. Upper limits of CN detections are indicated by an arrow inside the corresponding data point.
Because there were few CN detections in my data set, the N (CH)/N (CH+) ratio
was found to be more illustrative. The N (CH)/N (CH+) and N (CN)/N (CH+) ratios
provide similar qualitative measures of the gas density. Figure 4-13 displays a his-
togram of the N (CH)/N (CH+) ratio for my current sample along with similar data
from Pan et al. (2005) and Sheffer et al. (2008). My current sample is in hashed red
and the comparable sample is in gray. Above the histograms are the averages shown
as the symbols for each subset and with ±1 σ shown around each average. There is a vertical dashed line at N (CH)/N (CH+) = 1 indicating the distinction between CH+ like CH and CN like CH gas.
The categorization of CH+ like CH gas comes from the linear correlation between log[N (CH+)] and log[N (CH)] at lower column densities (log[N (CH+)] ≤ 12.6) where
Pan et al. (2005) find a slope of approximately one. At higher column densities
(log[N (CH+)] ≥ 12.6), there is only a very weak correlation if any between N (CH+)
and N (CH). However, at higher column densities, Pan et al. (2005) find that there is
a correlation between N (CN) and N (CN like CH), with slope of 1.7. The CN like CH
gas is what remains after accounting for the CH+ like CH gas at log[N (CH+)] = 12.6.
Hence, because there is a one to one ratio between CH+ and CH in CH+ like CH gas,
a N (CH)/N (CH+) ratio greater than one comes from the additional contribution of
179 CN like CH gas to the total amount of CH gas.
Figure 4-13: This figure shows the N (CH)/N (CH+) ratio histogram for the current sample against a combined comparable sample from Pan et al. (2005) and Sheffer et al. (2008). The current sample is in hashed red and the comparable sample is in gray. Above the histogram are symbols indicating the average value for each subset with ±1 σ range shown around each average. There is a vertical dashed line at N (CH) = N (CH+) indicating the distinction between CH+ like CH and CN like CH gas.
180 In addition, Figure 4-14 highlights the changes in density seen in Figure 4-12.
Figure 4-14 shows the N (CN)/N (CH+) ratio on a plot of the observed N (CO) versus
N (H2) after subtracting the N (CO) expected from the observed N (H2) based on the trends in N (CO) versus N (H2) plot. After subtracting the trend in N (CO) versus
N (H2), there are still clear density gradients. We see density increasing with the column density for both CO and H2. In the models by Sheffer et al. (2008) using Cloudy, they find that the sight lines corresponding to the upper half of this plot come from higher density or a lower UV radiation field, while sight lines in the bottom half typically have lower densities or higher UV fields.
Figure 4-14: This figure shows the variation in the N (CN)/N (CH+) ratio on the N (CO) versus N (H2) plot after accounting for the trends in N (CO) versus N (H2). The current sample is in red and the comparison sample is in black. Smaller symbols indicate lower densities and larger symbols indicate higher densities. See Figure 4-12 for further details on the sizes.
181 Table 4.11: CO and H2 Results for h and χ Percei
a −3 Sight Line log(N CO) T CO(1-0) log(N H2 ) T ex(1-0) MIN log(n /cm ) MAX HD 13841 14.72(0.04) 5.9(1.2) 20.60(0.10) 60(7) 2.50 2.70 2.85 HD 13969 15.15(0.06) 3.8(0.8) 20.87(0.10) 58(8) 1.60 2.20 2.50 HD 14053 14.51(0.09) 5.0(2.0) 20.96(0.10) 61(7) 1.55 2.55 2.80 BD+56◦ 0501 14.57(0.05) 4.0(0.6) 20.75(0.10) 57(10) 2.00 2.30 2.40 BD+56◦ 0563 14.38(0.09) 5.9(2.3) 20.62(0.10) 66(8) 2.10 2.70 2.90 HD 14443 14.36(0.08) 5.5(2.1) 20.98(0.10) 48(4) 2.00 2.65 2.95 HD 14476 14.91(0.09) 6.1(2.8) 20.86(0.14) 55(9) 2.00 2.70 2.95 BD+56◦ 0578 14.40(0.03) 6.4(1.2) 20.66(0.14) 66(11) 2.55 2.70 2.85 HD 14947 15.49(0.14) 3.0(1.0) 20.80(0.10) 55(6) ... 1.60 2.35 182 HD 15629 15.17(0.10) 5.0(2.0) 20.92(0.10) 67(11) 1.55 2.50 2.80 HD 16691 15.35(0.04) 3.1(0.4) 20.81(0.11) 48(4) ... 1.85 2.15 HD 17505 15.27(0.04) 3.3(0.4) 20.79(0.10) 66(8) 1.30 1.90 2.20 HD 17520 14.69(0.05) 5.7(1.5) 20.89(0.13) 48(8) 2.45 2.70 2.85 BD+60◦ 0586 14.89(0.03) 4.8(0.6) 20.72(0.10) 47(4) 2.45 2.55 2.70 HD 19243 14.66(0.04) 6.0(1.0) 20.84(0.10) 66(9) 2.55 2.65 2.75 HD 19820 15.48(0.06) 3.7(0.7) 20.71(0.10) 61(8) 1.60 2.15 2.50 Sheffer et al. (2008) HD 12323 14.53 3.1 20.32 82 ... 1.75 ... HD 13268 14.2 3.4 20.51 92 ... 1.90 ... HD 13745 13.94 4.0 20.67 66 ... 2.30 ... HD 14434 14.36 4.4 20.43 99 ... 2.35 ... HD 15137 13.52 3.1 20.32 104 ... 1.60 ... a Refers to the density of collisional species; (H+H2). 4.4.3 Discussion
The density results for the gas toward h and χ Per differ from the rest of the current sample and bias the overall sample. Compared to the locations of Sheffer et al.’s h and χ Per sample in the plot of CO versus H2, my sample of h and χ Per sight lines are at higher column densities, which should correspond to higher gas densities.
However, the sizes of the CN/CH+ ratio in Figure 4-12 indicate that the h and χ subsample have similar densities to Sheffer et al.’s h and χ Per sample.
The histograms of densities from CN chemistry, Figure 4-8, indicate similar den- sities between my sample and the data from Sheffer et al. (2008), within one sigma.
This is due to the fact that the CN and CN like CH gas is coming from denser gas. All the CN measurements in the h and χ Per sample are upper limits and are not biasing the result. This can also be seen in the histogram of CH/CH+, where CH is used instead of CN because of the few CN detections. In this figure it can be seen that the current sample is composed primarily of CH+ like CH gas. This is also true of the lower density sight lines from the Sheffer et al. (2008) sample without CN detections.
CO formation is less efficient for CH+ like CH gas than for CN like CH gas.
In the h and χ Per sample, most of the CO is coming from the CH+ like CH gas, while the Sheffer et al. (2008) sample has most of the gas coming from CN like
CH gas. Table 4.11 gives the density results for the h and χ Per subsample from the Sheffer et al. (2008) sample as well as the current sample. The histograms of densities from CO(T (2-1)) indicate a slightly higher density than the CO(T (1-0)) and are within one sigma of Sheffer et al. (2008) sample, unfortunately this is based on only a few J = 2 detections. However, several of the detections of J = 2 in 12CO are also associated with detections of 13CO. These are discussed in Section 4.5.
183 4.5 The 12CO/13CO Isotopic Ratio
The two isotopologues 12CO and 13CO are both found in the deeper regions of diffuse clouds. Deep inside dark molecular clouds we expect the 12CO and 13CO molecules to be self/mutually shielded and their ratio to be close to the ambient carbon isotope ratio. Millimeter-wave observations of dense molecular cores in the local ISM from Wilson et al. (1999) give an average of 69 ± 6. Observations of ratios with other proxies such as CN and CH+ confirm the ambient 12C/13C ratio in the local ISM to be about 70, as discussed in the work of Sheffer et al. (2007) and Ritchey et al. (2011). CO fractionation away from the ambient 12C/13C ratio comes from two competing effects, isotopic charge exchange and selective isotopic photodissociation.
Each process dominates under different physical conditions and are discussed below.
The histogram of the H2 temperatures shown in Figure 4-2 highlights that at low kinetic temperatures, which is indicative of higher density gas associated with higher
12CO column densities, we consistently detect 13CO. Sheffer et al. (2007) attribute part of the connection between the kinetic temperature and detections of 13CO to the exponential dependence of the isotopic charge exchange (ICE) reaction,
13C+ +12 CO →12 C+ +13 CO + 35 K.
The ICE process, when not in competition with other processes like selective pho- todissociation (SPD) that increases 12CO through self shielding, has an equilibrium isotopic ratio that depends only on the kinetic temperature (Sheffer et al. 2007),
−35K 12CO kr Tkinetic 13CO = 70 kf = 70 e .
Here, k f and k r refer to the forward and reverse rate coefficient for the ICE reaction
+ 12C with C and 70 is the ambient carbon isotope ratio, 13C = 70. In colder gas, the ICE
184 reaction is faster and more efficient in converting 13C+ to 13CO (Sheffer et al. 2007).
4.5.1 Regimes
Table 4.12 presents the column density, excitation and isotopic ratios for all sight lines with detectable amounts of 13CO for my data set. Figure 4-15 shows the isotopic column density ratio of CO, N (12CO)/N (13CO), against the column density of 12CO.
The general locations of the different regimes defined by fractionation are indicated with labels in the figure. This figure is comparable to Figure 7 from Sheffer et al.
(2007).
There are four regions in Figure 4-15 that help show when isotopic charge ex- change (ICE) or selective isotope photodissociation (SPD) dominates. First is the region of high N (12CO)/N (13CO) (N (12CO)/N (13CO) ≥ 100) where 12CO lines are optically thick enough to be self shielded but 13CO lines are not, leading to enhanced photodissociation of 13CO. This is where selective SPD drives up the ratio by self shielding for the dominant isotopologue(s). Most of the sight lines in this region have detections of J = 2 and 3 lines from 12CO. The second region of interest in Figure
4-15 is where N (12CO) ≥ 6x1015cm−2. This region with high N (12CO) is where 12CO is optically thick enough to be self shielded and 13CO has become self shielded and mutually shielded by 12CO. As a result, the ratio is close to the ambient carbon iso- tope ratio (12C/13C) of 70. Most of the sight lines in this region also have detections of J = 2 and 3 lines in 12CO. There is one sight line from the current sample, HD
170740, which has the largest N (12CO) and the only detection of a J = 3 line of 12CO in our sample.
185 Table 4.12: 13CO Results
Sight Line NJ=0 NJ=1 NJ=2 NJ=3 Total NCO T(0−1) T(0−2) T(0−3) (1014 cm−2) (1014 cm−2) (1014 cm−2) (1014 cm−2) (1014 cm−2) (K) (K) (K) HD 108 (12CO) 5.0(0.1) 3.8(0.1) 0.46(0.12) ... 9.3(0.2) 4.0(0.1) 4.2(0.3) ... HD 108 (13CO) ≤ 0.6 ≤ 1.2 ...... ≤ 1.8 ...... HD 108 (12CO/13CO)a ...... ≥ 5.0 ...... HD 46223 (12CO) 3.4(0.1) 1.8(0.1) 0.21(0.01) ... 5.4(0.1) 3.2(0.2) 3.8(0.1) ... HD 46223 (13CO) 0.09(0.01) 0.06(0.02) ...... 0.15(0.02) 3.7(0.9) ...... HD 46223 (12CO/13CO)a ...... 34.7(4.6) ...... HD 108927 (12CO) 9.0(0.4) 9.3(0.5) 1.10(0.19) ... 19.4(0.6) 5.2(0.3) 4.5(0.2) ... 13 b b b
186 HD 108927 ( CO) 0.24(0.08) 0.25(0.17) ...... 0.49(0.19) 5.2(3.8) ...... HD 108927 (12CO/13CO)a ...... 37.3(14.4) ...... HD 170740 (12CO) 41.0(1.8) 58.0(2.8) 13.00(1.66) 1.6(0.3) 113.6(3.7) 7.4(0.6) 6.0(0.3) 6.4(0.3) HD 170740 (13CO) 0.59(0.14) 0.89(0.27) ...... 1.48(0.30) 8.0(4.4) ...... HD 170740 (12CO/13CO)a ...... 66.9(13.9) ...... HD 165783 (12CO) 1.7(0.1) 0.9(0.1) 0.08(0.02) ... 2.6(0.1) 3.2(0.1) 3.5(0.2) ... HD 165783 (13CO) 0.04(0.01) ≤ 0.05 ...... ≤ 0.09 ...... HD 165783 (12CO/13CO)a ...... ≥ 28.5 ...... HD 174509 (12CO) 19.1(0.3) 12.1(0.1) 0.94(0.03) 0.06(0.02) 32.2(0.7) 3.6(0.1) 3.6(0.1) 4.3(0.2) HD 174509 (13CO) 0.39(0.03) 0.17(0.02) ...... 0.56(0.05) 2.9(0.2) ...... HD 174509 (12CO/13CO)a ...... 51.5(4.8) ...... a The 12CO/13CO ratio is determined from the J = 0 and J = 1 combined column densities. b The J = 1 column density for 13CO is inferred from the fitted J = 0 column density and the 12CO excitation temperature. The third region of interest is between 70 ≤ (N(12CO)/N(13CO)) ≤ 100 and
N (12CO) ≤ 6x1015 cm−2. In this region neither 12CO and 13CO are shielded signif-
icantly and their column density ratio lies roughly about the ambient value of 70.
There are two sight lines from the GOT C+ sample that fall into this region, HD
165783 and HD 174509. The fourth region (12CO/13CO ≤ 70) is where isotopic charge
exchange dominates. The low 12CO/13CO ratio in this region comes from the ICE
reaction with C+, described by Watson et al. (1976), which converts 12CO into 13CO
through
12CO + 13C+ → 13CO + 12C+.
This region seems to be separated into two subregions of high and low 12CO column
density, N (12CO) ≤ or ≥ 6x1015 cm−2. This separation is supported by the additional
millimeter-wave observations from Liszt & Lucas (1998), which are found within the
ICE dominated domain at both low and high 12CO column density. There are two sight lines from the current data set that fall into the ICE dominated regime, HD
46223 and HD 108927. These are discussed further in the next section.
4.5.2 Density Histograms for 12CO/13CO Sample
Figure 4-16 shows a histogram of the density of collisional partners derived from the CO(J = 1 to J = 0) temperature for the sight lines presented in Figure 4-15. The
12 13 sight lines with shielded CO and shielded CO have a higher average density (nH+H2 ≈ 375 cm−3). These sight lines are referred to as the shielded sample in Figure 4-16.
The sight lines with only 12CO shielded (Region 1) and those with unshielded 12CO
−3 (Region 3) have lower densities (nH+H2 ≈ 130 cm ) that fall into roughly the same region within uncertainties. The sight lines with low 12CO/13CO ratios controlled by
ICE can be divided into two groups, low N (12CO) and high N (12CO). Low N (12CO)
−3 12 ICE sight lines have the lowest densities (nH+H2 ≈ 100 cm ), while high N ( CO)
187 200
Sheffer et al. (2007) Sample 150 Listz & Lucas (1998) Sample Current Sample
100
50
ICE
0 0 0.5 1 1.5 2
Figure 4-15: This figure shows the isotopic column density ratio of CO against the column density of 12CO for my sight lines and pre- vious values. Black filled squares correspond to previous values from UV measurements (Sheffer et al. 2007). Black open circles correspond to values from millimeter-wave observations of Liszt & Lucas (1998). Red triangles are the detections and upper lim- its from my FUSE sample and blue triangles are from the GOT C+ sample. The general locations of the different regimes are indicated with labels. This figure has been adapted from Sheffer et al. (2007).
188 −3 ICE sight lines have the highest densities in this small sample (nH+H2 ≈ 550 cm ). The new observations of 13CO are included in the histogram of Figure 4-16. The gas toward HD 46223 is consistent with the low N (12CO) ICE group of sight lines. The
gas toward HD 108927 seems to fall on the border between the high and low N (12CO)
ICE regimes in Figure 4-15. A plot of the isotopic ratio versus 13CO (Figure 4-17)
helps show that HD 108927 should fall into the high N (12CO) ICE category, because
its location on the plot is separate from the low N (12CO) ICE sight lines. The density
inferred from 12CO(1-0) toward HD 108927 is distinct from the low N (12CO) ICE
group of densities. The density inferred from 12CO(2-1) has error bars that extend
into both categories. The diffuse molecular gas toward HD 170740 is consistent with
the high density shielded 12CO and 13CO group; this is based on both the location in
the plot of Figure 4-15 and on the density inferred from the 12CO(1-0) and 12CO(2-1)
analyses.
Figure 4-18 shows a histogram of the density derived from the CO(J = 2 to J = 1)
temperature. There is more uncertainty here than in Figure 4-16, but the results are
similar. The unshielded and the majority of the 12CO shielded sight lines have lower
−3 −3 12 density (nH+H2 ≈150 cm and 125 cm , respectively), while the shielded CO and
13 −3 CO sight lines have higher density (nH+H2 ≈375 cm ). The ICE sight lines still fall into two subgroups; however, they are closer in density than the densities inferred
from CO(1-0). An analogous histogram of the density of collisional partners derived
from the CN chemistry is not shown because the low number of CN detections result in
the uncertainties of the different regimes overlapping and becoming indistinguishable.
189 Figure 4-16: This figure shows a histogram of density derived from the J = 1 to J = 0 excitation temperature of 12CO. The sight lines in this figure include the same data sets as in Figure 4-15 but are grouped by the different regimes shown in Figure 4-15 and described in the text.
190 200
Sheffer et al. (2007) Sample 150 Current Sample
100
50
ICE
0 0 0.01 0.02 0.03 0.04 0.05
Figure 4-17: This figure shows the isotopic column density ratio of CO against the column density of 13CO for my sight lines and results from previous analyses. Black filled squares correspond to val- ues from previous UV measurements (Sheffer et al. 2007). Red triangles are the detections and upper limits from my FUSE sample and blue triangles are from the GOT C+ sample. The general locations of the different regimes are indicated with la- bels. The millimeter-wave observations of Liszt & Lucas (1998) have N (13CO) values beyond the range shown here.
191 Figure 4-18: This figure shows a histogram of density derived from the J = 2 to J = 1 excitation temperature of 12CO. The sight lines in this figure include the same data sets as in Figure 4-15 that have J = 2 12CO detections. The data are grouped by the regimes shown in Figure 4-15 and discussed in the text.
192 4.6 Kinematic Associations of GOT C+ Species
To illustrate the general results, a component-by-component analysis is shown for
G014.8-1.0 (see Table 4.13). The remainder of the sight lines are discussed in terms of connections to nearby dark clouds and notable deviations from the general results.
Upper limits for the linear separation are calculated using the derived distance to the target star and the observed angular separation; if two targets are near the same pointing, the distance to the nearest target is used in the calculations.
4.6.1 A Component by Component Discussion
for G014.8-1.0
HD 168607 is at a distance of 1100 pc and is 12 arcmin away (≤ 3.8 pc at 1100 pc) from the pointing G014.8-1.0. Almost all the components seen in H I are also seen in Ca II, indicating they are both probes of diffuse regions of the individual clouds.
However, there are many components seen in Ca II that are not seen in emission
(at Vlsr of -45.5, -39.7, -33.6, -28.4, -18.2, -13.4, -5.3, -0.8, 3.4, 9.9, 59.2, 65.6 km s−1). This difference is due to the increased sensitivity of absorption measurements relative to emission measurements. There are also some H I features that do not have corresponding absorption features; these features have large Vlsr (70.0 and 112.6 km s−1) and are most likely probing the material behind HD 168607.
+ There are three components that are seen in Ca II, CH , and CH (near Vlsr of 6.7, 15.0, 23.3 km s−1); these components only correspond with emission features seen in H I. The presence of CH and CH+ indicates that these components are probing denser regions than if only Ca II or H I were observed. We would expect these types of regions, with CH or CH+ detections, to have some amount of CN or CO. The
CH+ absorption indicates the presence of relatively low density gas (Pan et al. 2005);
+ otherwise, reactions with H2 would destroy CH . There are clear relationships involv-
193 19 −2 ing CO and CH with molecular hydrogen for N (H2) ≥ 10 cm , as documented in
Sheffer et al. (2008). These H2 column densities correspond to a minimum molecular hydrogen fraction of 5% to 10%, where the molecular hydrogen fraction is defined as
2N (H2)/(2N (H2)+N (H)). This is the minimum molecular hydrogen fraction we use to define gas as molecular. The consequences of CO production were noted in Section
4.3. As a reminder, in low density gas, the CH+ + O → CO+ + H reaction leads
to CO production. The CH and CO are also coupled through the neutral-neutral
reaction route CH + O → CO + H. However, the abundances of these species are not
high enough for a reliable detection in CO emission. Therefore, these components
are most likely probing regions of CO dark gas, where CO is present but at a column
density too low to be confidently measured in emission.
There is one component with detectable CN, CH, and Ca II near a Vlsr of 18.7 km s−1. This component corresponds to emission in C+, 12CO, 13CO, and C18O. The
presence of 12CO, 13CO, and C18O all indicate that this component is probing much
denser gas. Here, the associated data from absorption are revealing the diffuse outer
regions of a molecular cloud.
There are several components seen in CO emission that are not seen as molecular
−1 absorption components (near Vlsr of -26.1, 30.0, 39.1, 46.3, 55.9, 75.6 km s ), de-
spite the majority being associated with components in Ca II. The presence of CO
emission indicates that the density in the area being probed should be high enough
for detectable amounts of other species, such as CH or CH+. There will be some
differences that arise from variations in the small scale structure between the differ-
ent locations of the emission and absorption measurements, in this case 12 arcmin.
Thus, the presence of only Ca II indicates that the absorption measurements are likely
probing the lower density atomic region that surrounds the molecular gas associated
with the CO emission. Table 4.13 provides a summarized component-by-component
interpretation for this sight line. The stacked spectra are shown in Figure E-1.
194 Table 4.13: Comparison of Component Structure: G014.8-1.0
G014.8-1.0 HD 168607 Vlsr GOT C+ McDonald Comments (km s−1) Components Components -24 H I, 12CO Ca II Molecular cloud edge -9.0 H I Ca II Warm neutral 6.0 H I Ca II, CH+, CH Cold neutral, CO dark 15 H I Ca II, CH+, CH Cold neutral, CO dark 19 C+, 12CO, 13CO, C18O Ca II, CH, CN Molecular cloud 23 H I Ca II, CH+, CH Cold neutral, CO dark 30 12CO, 13CO, C18O Ca II Molecular cloud 37 H I Ca II Warm neutral 46 12CO Ca II Molecular cloud edge 53 H I, 12CO Ca II Molecular cloud edge
4.6.2 Summaries for Other Pointings
Tables 3.20 and 3.21 show the Vlsr structure of components with emission features and the Vlsr of any corresponding absorption features for the inner and outer Galaxy. However, there are often components seen in absorption that do not correspond with any of the components seen in emission. These components are discussed here.
G010.4+0.0: HD 165783 is 4 arcmin away and HD 165918 is 6 arcmin away from
G010.4+0.0. Assuming a distance of 500 pc, the separations between directions re- lated to the absorption and emission measurements are ≤ 0.6 pc and ≤ 0.9 pc, respectively.
There are five components seen in emission and toward both HD 165783 and HD
−1 165918 (near Vlsr of -21.2, -10.3, 0.0, 4.2, 10.7 km s ). Three of these are only associated with H I emission and are probing lower density regions (near Vlsr of -21.2,
0.0, 10.7 km s−1). One component, seen in Ca II at 4.2 km s−1, is also associated with
CH/CH+ absorption and C+ emission. The presence of CH and CH+ without CO indicates this is likely probing diffuse molecular gas and regions of CO dark gas.
195 Table 4.14: Comparison of Component Structure: G010.4+0.0
G010.4+0.0 HD 165783 HD 165918 Vlsr GOT C+ McDonald McDonald Comments (km s−1) Components Components Components -20 H I Ca II Ca II Warm neutral -10 H I, 12CO Ca II, CH+ Ca II, CH+ Molecular cloud edge 1.0 H I Ca II Ca II Warm neutral 5.0 C+ Ca II, CH+, CH Ca II, CH+, CH Cold neutral, CO dark 7.0 12CO Ca II, CH+, CH ... Molecular cloud edge 10 H I Ca II Ca II Warm neutral 25 C+, H I, 12CO Ca II ... Molecular cloud edge 35 H I, 12CO Ca II ... Molecular cloud edge 45 C+, H I Ca II ... Cold neutral
There are six components from species seen in absorption without corresponding
emission features. These additional components are seen toward HD 165783 near Vlsr
(km s−1) of -15.9 (Ca II and CH+), -6.4 (Ca II), -3.6 (Ca II and CH+), and 13.5,
16.6, 39.0 (Ca II). Only three of these additional components are seen toward HD
−1 165918 near Vlsr (km s ) of -16.0, -3.5, 16.0 (Ca II). These differences likely stem from small scale variations in structure and weak Ca II components. There are several nearby dark clouds including ZHS2005 C3C, C3B, C3A from Zhang et al. (2005) but all measured Vlsr are beyond the observed range spanned by the sight lines toward
18 HD 165783 or HD 165918. The Vlsr given by Zhang et al. (2005) for C O at 66 km s−1 agrees with our measurements of a component seen in C18O, C+, and H I at 65 km s−1. Table 4.14 shows the summarized component-by-component interpretation for this GOT C+ pointing for both of its two nearby stellar sight lines. See Figures
3-17 and E-2 for the stacked spectra.
G015.7+1.0: HD 167498 is 9 arcmin away and HD 167812 is 15 arcmin away.
Assuming a distance of 250 pc, the separations between the directions observed in absorption and emission are ≤ 0.7 pc and ≤ 1.1 pc, respectively. There is one com-
196 ponent seen in emission from G015.7+1.0 with H I and toward both HD 167498 and
−1 HD 167812 in Ca II and CH (near a Vlsr of ∼4.1 km s ). This component is likely probing CO dark gas. There are two weak Ca II components toward HD 167498 that have no emission features associated with them (near Vlsr of -14.7 and -10.8 km s−1). There is also one component toward HD 167498 and HD 167812 observed in
+ absorption from CH, CH , and Ca II that has no emission feature (near a Vlsr of 1.0 km s−1). This component is also likely probing CO dark gas. The stacked spectra are shown in Figures E-3 and E-4.
G020.0+0.0: HD 169754 is 14 arcmin away from G020+0.0. Assuming a distance of
2000 pc, the directions are ≤ 8.1 pc apart. Almost all the H I emission components are also seen in Ca II. There are two components from species detected via absorption toward HD 169754 that are unassociated with components seen in emission near Vlsr
(km s−1) of -7.3 (Ca II) and 12.2 (Ca II and CH+). There are several nearby dark clouds including L413, from Lynds (1962), but only one with additional velocity information. BGPS G020.050+0.195 has a velocity component around 18 km s−1
+ + seen in HCO , N2H , and NH3 from Schlingman et al. (2011) that agrees with a component seen in Ca II,H I and CH+ near 20 km s−1, and is likely associated with the outer envelope of the molecular cloud. The stacked spectra are shown in Figure
E-5.
G032.6+0.0: HD 174509 is 13 arcmin away from G032.6+0.0 and assuming a dis- tance of 500 pc, the directions probe material ≤ 1.9 pc apart. The component with
H I emission at -0.5 km s−1 is associated with Ca II, Ca I,K I, CH+ and CH absorp- tion and is likely probing diffuse molecular gas. The component with H I and C+ emission near 14.4 km s−1 is associated with Ca II and Ca I indicating diffuse atomic gas. The other three absorption components all have detections of Ca II, Ca I,K I,
CH and CN; these are associated with CO emission near Vlsr of 5.8, 9.7, and 12.1 km
197 s−1 and are likely probing the cloud envelope.
There are several dark clouds nearby with measurements in NH3: BGPS G032.524- 00.131 (79.3 km s−1), BGPS G032.527-00.123 (79.4 km s−1), BGPS G032.411-00.0099
(80.2 km s−1), and BGPS G032.436-00.17 (77.4 km s−1) from Dunham et al. (2011).
However, all the corresponding Vlsr measurements lie beyond the range measured for species in absorption. There is agreement with a component seen in emission in 12CO,
13CO, C18O, C+, and H I near 75 km s−1, indicating this molecular cloud lies behind
HD 174509. The stacked spectra are shown in Figure E-6.
G091.7+1.0: BD+49◦ 3482 is 9 arcmin away and BD+49◦ 3484 is 4 arcmin away
from G091.7+1.0. For a distance of 250 pc, the directions are ≤ 0.7 pc and ≤ 0.3 pc
apart, respectively. It should be noted that the pointing G091.7+1.0 did not include
CO measurements. There is one component seen in H I emission and in Ca II and
+ ◦ ◦ −1 CH for both BD+49 3482 and BD+49 3484 near Vlsr of 3.5 km s . There is also
a component near 1.6 km s−1 seen toward BD+49◦ 3482 that is associated with K I
and CH+ absorption and H I emission.
◦ −1 The weak Ca II component toward BD+49 3482 (at Vlsr of -4.4 km s ) has no
emission features associated with it, and three Ca II components toward BD+49◦ 3484
−1 also have no associated emission features (at Vlsr of 8.5, 10.3 and 13.0 km s ). Two of the components toward BD+49◦ 3484 also have a CH component (near 8.5 and
10.3 km s−1). One of these components also has CN absorption (near 10.3 km s−1).
There are two dust clouds along adjacent directions, Dobashi 2934 from Dobashi
(2011) and TGU H541 P30 from Dobashi et al. (2005). There is also a molecular
13 −1 cloud, [DBY 94] 091.7+00.9 observed in CO with Vlsr of 10.9 km s from Dobashi et al. (1994). This agrees nicely with the absorption seen in CN (10.2 km s−1), CH
(10.6 km s−1), and Ca II (10.3 km s−1) and C+ emission, indicating the presence of a diffuse molecular envelope. The nearby sight line (BD+49◦ 3482), although farther
198 from the Herschel pointing, mainly shows the lower velocity material. The stacked
spectra appear in Figures E-7 and E-8.
G109.8+0.0: HD 240179 and HD 240183 are both 7 arcmin away from G109.8+0.0
and have a separation of ≤ 1.4 pc, assuming a distance of 700 pc. It should be noted that the pointing G109.8+0.0 did not include CO measurements. There are two components seen in H I emission and Ca II absorption toward both HD 240179
−1 and HD 240183 (near Vlsr of -10.3 and 0.2 km s ). Spectra toward HD 240179 also reveal this component in CH+, as well as an additional component near 4.0 km s−1 that is not seen toward HD 240183. This component near 4.0 km s−1 is detected in
H I emission and Ca II and K I absorption and is likely sampling a diffuse molecular cloud edge.
There are several components seen in absorption that have no corresponding emis-
−1 sion features (near Vlsr of -7.0, -3.5, -1.0, 7.6, 11.9 and 17.5 km s ), several of which
−1 are only weak Ca II components (near Vlsr of 7.6, 11.9, and 17.5 km s ). However, there are three components with Ca II, CH, CH+, and K I with no corresponding
−1 emission features (near Vlsr of -7.0, -3.5, -1.0 km s ) and might point to the pres- ence of CO dark gas, but without CO emission data for this pointing it is difficult
to confirm. There is a dense core nearby, IRAS 23033+5951, which has been studied
extensively. NH3 observations by Wouterloot et al. (1988) show emission at Vlsr of -52.8 km s−1, while 12CO and 13CO measurements by Wouterloot & Brand (1989) in-
−1 dicate respective Vlsr of -51.8 and -53.1 km s . This component is clearly associated
with the emission seen in H I and C+. The stacked spectra are shown in Figures E-9
and E-10.
G207.2-1.0: HD 47073 and HD 260737 are both 7 arcmin away from G207.2-1.0,
and assuming a distance of 200 pc, the directions are both separated by ≤ 0.4 pc.
There is one component seen in H I emission and in Ca II absorption for both HD
199 −1 47073 and HD 260737 near a Vlsr of 3.8 km s and is probably only probing diffuse atomic gas. There is one additional component toward HD 260737 that is seen in
12 both emission (H I and CO) and absorption (Ca II and K I) near a Vlsr of 9.4 km s−1 and is likely probing the envelope of a molecular cloud. Four additional weak
Ca II components toward HD 260737 are not seen in emission (-15.6, -8.6, -2.9, 1.8, and 5.3 km s−1). The stacked spectra are given in Figures E-11 and E-12.
G225.3+0.0: HD 55469 is 26 arcmin away and HD 55981 is 8 arcmin away from
G225.3+0.0. Assuming a distance of 400 pc, the separations are ≤ 3.0 pc and ≤ 0.9 pc, respectively. There are no H I observations available for this pointing. Emission
+ 12 components are detected toward G225.3+0.0 near Vlsr of 17.8 (C and CO) and 51.9 (C+) km s−1. There are absorption components detected toward HD 55981 near
+ −1 Vlsr of -7.3, -2.3, 2.8, 5.7, 9.2 (Ca II), 14.8 (Ca II,K I, and CH ), and 20.9 km s
(Ca II). There are also absorption components detected toward HD 55469 near Vlsr of 2.0 (Ca II and K I) and 3.8, 10.3 km s−1 (Ca II). While there are components seen in emission and in absorption, no clear correspondence is found between the emission and absorption components. The closest correspondence is seen for the absorption
+ −1 + from CH , K I, and Ca II around a Vlsr of 14.8 km s and the emission from C
12 −1 13 and CO around a Vlsr of 17.5 km s . Kim et al. (2004) studied L1658 in CO at
−1 coordinates G224.3-1.1 (at a Vlsr of 15.5 km s ) and G226.1-0.04 (at a Vlsr of 16.2 km s−1). These velocities both fall between those of the emission at 17.5 km s−1 and the absorption at 14.8 km s−1, suggesting that the components are related. The stacked spectra are provided in Figures E-13 and E-14.
200 4.6.3 Updated Parallax Distances
Recently, Bailer-Jones et al. (2018) performed a reanalysis of the inferred distances from the parallax measurements in GAIA Data Release 2. This reanalysis led to more accurate distances; the parallax distances are listed in Tables 2.1 and 2.8 and have errors less than 20%. In the primary sample, there are eight sight lines with significant changes in distance. However, these changes do not impact our results from our primary sample.
For the GOT C+ sample, these changes affect the maximum projected distances between the absorption sight lines and emission pointings. The three sight lines with the largest differences in inferred distance are HD 55918, HD 165783, and HD 165918.
Fortunately, when there are multiple sight lines near a pointing, our calculations for the projected separation are based on the distance to the nearest of the two stars. This leads to only small changes in the separations. The largest change in the projected separation (an increase of 1.4 pc) was between HD 168607 and G014.8-1.0, all other changes in the separations were less than 1 pc. These revised distances do not change our final results.
201 Chapter 5
Discussion
There are a wide range of densities between atomic and molecular clouds and only a complete set of density tracers can create a full picture. Section 5.1 describes the combined picture of the different density tracers from our primary data set. Section
5.2 extends this picture by incorporating the results of the kinematic association from the GOT C+ project. Results from both cloud structure and density are applied to the analysis of several well studied regions in Section 5.3, where the prominence of
CO dark gas is highlighted.
5.1 Combined Chemistry Picture
The carbon budget in molecular clouds is determined from the balance between the production of CO from ion neutral pathways with some neutral neutral reactions and the destruction of CO from photodissociation and reactions between He+ and CO.
In the deeper regions of molecular clouds, photodissociation is drastically reduced due to the effects self shielding and mutual shielding with coincident lines from atomic and molecular hydrogen. As a result, almost all the carbon is able to remain in the form of CO and the small amounts of atomic carbon arise from its destruction.
In the more diffuse regions of molecular clouds (where the molecular hydrogen fraction is at least 5% to 10%), most of the carbon is singly ionized (C+) and little 202 or no CO is present, as CO will easily photodissociate into atomic carbon and then be photoionized into C+. In this region, there is an abundant number of electrons to destroy molecular ions through dissociative recombination, which reduces the effec- tiveness of the ion molecule chemistry that dominates in denser regions. Photoion- ization becomes less efficient as the depth into a cloud increases. In the translucent cloud regime often referred to as diffuse molecular gas, modeling efforts indicate that the primary form of carbon transitions to neutral carbon and then to CO. While C+ is still present, the predicted abundance has dropped by an order of magnitude. It is worth noting, however, observations of C I (e.g., Jenkins & Tripp 2011) and CO
(Sheffer et al. 2008) reveal column densities that are comparable or favor CO by up to factors of 10.
There are a number of additional carbon species that are important indicators of density. CN production proceeds from CH to C2 and then to CN, making CN a third generation neutral molecule as discussed in the work of Federman et al. (1984) and in Section 4.3.1. This means that CN is only produced after significant amounts of the precursor molecules, CH and C2, are available. Table 5.1 is a compilation of our results from different density probes. However, there is limited information for a number of species in our data set, and so we have also provided Table 5.2, which is a compilation of results of our analysis for previous measurements, with their references listed in the footnotes. This combined sample provides the basis for our interpretation. All densities shown in Tables 5.1 and 5.2 refer to the density of collisional partners.
203 Within Table 5.1 is the observed CN column density for each component or an
upper limit for each sight line along with the amount predicted from CN chemistry.
This CN column density arises from the chemical model with the density of collisional
partners listed in the next column; when there are multiple components of CN, an
average weighted by the fractional amount of CN in each component is also given for
comparison. The next set of columns comes from the CO excitation temperatures.
There are some detections of J = 2 but only one detection of J = 3, which is given
as a footnote (HD 170740). There are a few measurements of C2 column density that were used in the determinations from CN chemistry. The densities derived from
C2 are based on C2 excitation (e.g., Sonnentrucker et al. 2007, Hupe et al. 2012). There are only two determinations using C+ observations and are listed in the table.
The last columns show the density of collisional partners derived from neutral carbon excitation. There are sometimes additional components seen in neutral carbon that are not seen in CN; because of this, we list two columns under densities derived from C I. The range for the combined line of sight, using all the neutral carbon components, is given under the ‘all’ column and the density range from only the neutral carbon components with corresponding CN components are listed under the
‘CN Only’ column.
The compilation of results from the previous work, shown in Table 5.2, have been reanalyzed using the same methods and constants that were adopted in analyzing the current sample. The original sources are listed in the footnotes and are referenced throughout the table. Table 5.2 includes the same density measures shown in Table
5.1; however, there are more data for several tracers in this combined sample than for the current sample alone, helping to provide a more solid basis for our conclusions.
204 In general, the gas densities derived from C I and C2 excitation are among the low- est densities probed. These species probe similar volumes of gas that have only rela-
12 tively small amounts of CO present. Higher gas densities are probed with T CO(1-0),
+ C excitation, CN chemistry, and T CO(2-1). On average the densities derived from
CN chemistry fall between those from T CO(1-0) and T CO(2-1), but are typically closer
to the densities from T CO(1-0). The CO gas being probed by the CO(1-0) transition and the CN molecule are
known to occupy similar volumes deeper into the diffuse molecular cloud (Pan et al.
2004). The densities from T CO(2-1) are within a factor of two of the densities probed by CN, but are consistently higher, indicating that the more excited CO resides deeper
into the cloud. The densities derived from T CO(3-2) should probe the innermost regions of the cloud; however, there are too few detections to draw strong conclusions.
C+ is found throughout the cloud and since the elemental abundance of carbon
should not vary throughout the cloud (Sofia et al. 1997), one expects a lower density
than that found from CN chemistry, which is from a denser region deeper into the
cloud. However, C+ emission is sensitive to pressure and the density of the gas and
arises preferentially from denser regions. From the average line of sight densities
derived from C+ excitation, it appears that C+ emission from diffuse molecular gas
resides between where the gas probed by CN chemistry and T CO(1-0) are located within the cloud. In Section 5.3, these results are combined with those from the
GOT C+ analysis given in Section 5.2.
205 Table 5.1: Comparison of Density Determinations (Current Data Set)
Sight Line V lsr N (CN) nH+H2 nH+H2 (TCO) nH+H2 nH+H2 nH+H2 (C I) Observed Predicted (CN) (J = 1-0) (J = 2-1) (C2) (C II) (all) (CN Only) (km s−1) (1011 cm−2) (1011 cm−2) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) HD 108a -5.1 15.0 15.0 400 ...... 6.9 11.0 11.0 975 ...... Combinedb ...... 650 200 325 ...... 25-90 35-100 HD 5689 -12.6 ≤ 44.0 ≤ 44.3 ≤ 450 100 ......
206 HD 12882 0.2 ≤ 27.0 ≤ 27.1 ≤ 425 150 ...... HD 13841 -41.6 ≤ 7.0 ≤ 7.0 ≤ 175 500 ...... 10-150 5-200 HD 13969 -10.8 ≤ 16.0 ≤ 15.6 ≤ 300 150 ...... HD 14053 -8.6 ≤ 23.0 ≤ 22.9 ≤ 275 350 ...... BD+56◦ 0501 -9.2 ≤ 28.0 ≤ 27.8 ≤ 600 200 ...... BD+56◦ 0563 -8.4 ≤ 15.0 ≤ 14.9 ≤ 225 500 ...... HD 14443 -6.2 ≤ 12.0 ≤ 11.7 ≤ 225 450 ...... HD 14476 -0.6 ≤ 22.0 ≤ 22.0 ≤ 200 500 ...... BD+56◦ 0578 -12.4 ≤ 32.0 ≤ 32.3 ≤ 675 500 ...... HD 14947 -7.7 ≤ 17.0 ≤ 16.7 ≤ 75 50 ...... HD 15629 -5.1 ≤ 36.0 ≤ 35.7 ≤ 450 325 ...... HD 16691 -12.7 ≤ 83.0 ≤ 82.8 ≤ 1075 75 ...... Table 5.1: —Continued
Sight Line V lsr N (CN) nH+H2 nH+H2 (TCO) nH+H2 nH+H2 nH+H2 (C I) Observed Predicted (CN) (J = 1-0) (J = 2-1) (C2) (C II) (all) (CN Only) (km s−1) (1011 cm−2) (1011 cm−2) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) HD 17505 -4.3 ≤ 14.0 ≤ 14.3 ≤ 55 75 ...... HD 17520 -2.5 ≤ 22.0 ≤ 22.0 ≤ 375 500 ...... BD+60◦ 0586 -2.2 ≤ 17.0 ≤ 17.0 ≤ 225 350 500 ...... HD 19243 0.5 ≤ 8.0 ≤ 7.7 ≤ 175 450 ...... HD 19820 -14.9 ≤ 25.0 ≤ 25.0 ≤ 65 150 ......
207 HD 25443 -8.1 ≤ 16.0 ≤ 15.5 ≤ 150 ≤ 100 ...... 10-35 5-25 HD 25638 2.4 38.0 39.2 225 ...... 5.7 20.0 20.9 100 ...... Combinedb ...... 175 75 ...... HD 45314 1.7 ≤ 18.0 ≤ 17.7 ≤ 350 200 200 ...... HD 46223 11.3 ≤ 13.0 ≤ 13.0 ≤ 125 100 75 ...... 5-25 5-10 HD 92964 -9.8 0.5 0.4 50 ...... -6.2 2.6 2.0 75 ...... -0.2 3.9 4.1 125 ...... Combinedb ...... 100 575 ...... HD 97253 -14.6 32.9 38.7 225 ...... -7.3 1.6 1.3 75 ...... Combinedb ...... 225 325 . . . 85 ...... Table 5.1: —Continued
Sight Line V lsr N (CN) nH+H2 nH+H2 (TCO) nH+H2 nH+H2 nH+H2 (C I) Observed Predicted (CN) (J = 1-0) (J = 2-1) (C2) (C II) (all) (CN Only) (km s−1) (1011 cm−2) (1011 cm−2) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) (cm−3) HD 108927 1.5 20.0 21.7 250 400 325 85 . . . 20-40 20-40 HD 149404 -11.8 6.8 7.0 125 ...... 7.4 36.1 35.1 275 ...... Combinedb ...... 250 450 700 ...... HD 151932 0.1 0.8 0.8 50 ...... 6.6 26.2 31.8 175 ...... Combinedb ...... 175 350 625 85 ...... 208 HD 168076 0.9 17.5 21.2 225 ...... 6.5 3.3 3.9 175 ...... Combinedb ...... 225 275 . . . 125 920 ...... HD 170740c 6.6 78.3 90.7 450 700 800 85 480 15-95 15-100 HD 192303 8.9 ≤ 38.0 ≤ 38.1 ≤ 750 575 ...... HD 192641 4.0 ≤ 15.0 ≤ 15.2 ≤ 100 625 ...... HD 193793 6.2 ≤ 7.0 ≤ 7.1 ≤ 25 275 400 ...... HD 197770 -1.4 70.0 82.6 150 125 100 125 ...... HD 214419 -19.5 56.0 56.3 750 50 ...... HD 217086 -15.7 ≤ 25.0 ≤ 26.0 ≤ 25 200 ...... aHD 108 has an additional density determination from inferred from C I pressure measurements, 130 cm−3. bThe results for each of the CN components were weighted by their respective CN column densities and averaged for comparison with line of sight values from other measures. cHD 170740 has an additional density determination from CO excitation temperature (J = 3-2), 400 cm−3. Table 5.2: Comparison of Different Density Determinations (previous efforts noted in the text)
−3 a b Sight Line nH+H2 (T CO) (cm ) nH+H2 (CN) nH+H2 (C2) nH+H2 (C II) nH+H2 (C I) T H2 T C2 (J =1-0) (J =2-1) (J =3-2) (cm−3) (cm−3) (cm−3) (cm−3) (K) (K) BD+48◦ 3437 ≤10 ...... 1752 ...... 831 ... BD+53◦ 2820 80 ...... ≤6002 ...... 931 ... HD 12323 60 650 ... 5002 ...... 821 ... HD 13268 80 ...... 3502 ...... 921 ... HD 13745 200 ...... 2752 ...... 661 ... HD 14434 220 ...... 702 ...... 991 ... 2 1
209 HD 15137 40 500 ... ≤250 ...... 104 ... HD 22951 ≤10 ...... 802 ...... 401 ... HD 23180 140 ...... 1502 ...... 401 ... HD 23478 130 200 50 1752 505 ...... 551 20-405 HD 24190 60 ...... 1002 ...... 661 ... HD 24398 130 300 150 2502 756 ...... 302 50c6 HD 24534 320 900 160 1252 1006 3207 5107 541 456 HD 27778 450 1000 150 2502 1006 ...... 511 506 HD 30122 180 250 ... 1502 ...... 611 ... HD 36841 ≤10 ...... 1752 ...... 652 ... HD 37367 70 ...... ≤1252 ...... 731 ... HD 37903 ≤10 ...... 170007 5957 641 ... HD 43818 260 ...... ≤2252 ...... 652 ... HD 58510 20 ...... ≤1252 ...... 901 ... Table 5.2: —Continued
−3 a b Sight Line nH+H2 (T CO) (cm ) nH+H2 (CN) nH+H2 (C2) nH+H2 (C II) nH+H2 (C I) T H2 T C2 (J =1-0) (J =2-1) (J =3-2) (cm−3) (cm−3) (cm−3) (cm−3) (K) (K) HD 63005 130 ...... 4252 ...... 781 ... HD 94454 160 ...... 2007 557 741 ... HD 96675 160 2500 ... 4002 ... 2607 ... 502 ... HD 99872 160 200 ... ≤702 ...... 661 ... HD 102065 140 ...... ≤7502 ...... 652 ... HD 115071 140 ...... 2807 607 711 ... 210 HD 115455 30 ...... ≤1252 ...... 811 ... HD 137595 180 450 ... ≤1252 ...... 721 ... HD 144965 200 900 ...... 807 857 701 ... HD 147683 360 1800 900 ... 255 ... 1357 581 50-605 HD 147888 1260 2500 2500 150d4 1006 12607 2207 441 556 HD 147933 ≤10 2000 2000 2254 ...... 504 ... HD 148184 140 700 800 702 756 ...... 602 506 HD 149757 100 320 110 1752 756 ...... 602 506 HD 152590 220 ...... 9007 807 641 ... HD 154368 40 1000 ... 2252 756 4007 ... 471 206 HD 157857 250 ...... ≤4252 ...... 861 ... HD 177989 140 200 ...... 555 ...... 491 205 Table 5.2: —Continued
−3 a b Sight Line nH+H2 (T CO) (cm ) nH+H2 (CN) nH+H2 (C2) nH+H2 (C II) nH+H2 (C I) T H2 T C2 (J =1-0) (J =2-1) (J =3-2) (cm−3) (cm−3) (cm−3) (cm−3) (K) (K) HD 185418 100 200 ... ≤202 ... 1007 257 652 ... HD 190918 ≤10 500 ... ≤3502 ...... 1021 ... HD 192035 70 300 ... 450e2 955 ...... 681 20-305 HD 192639 ≤10 ...... ≤252 ... 2607 507 652 ... HD 198781 100 200 ... 2752 ≤255 ...... 651 10-1005 HD 203532 500 500 ...... 555 ...... 471 20-405 HD 206267 500 1000 360 2504 1256 5207 657 653 356 HD 206267 ...... 2504 755 ...... 20-405 211 HD 207198 160 160 ... 1504 1006 2607 657 663 606 HD 208905 450 ...... ≤2752 ...... 771 ... HD 209481 30 ...... 504 ...... 781 ... HD 209975 30 ...... ≤2002 ...... 731 ... HD 210121 560 1600 30 225f 2 1006 ...... 511 456 HD 210809 50 ...... ≤3002 ...... 871 ... HD 210839 200 320 ... 2254 506 1407 657 504 30c6 HD 220057 40 300 ... 2502 455 ...... 651 905 Data sources: 1 – Sheffer et al. (2008) Table 6; 2 – Sheffer et al. (2008) Table 7; 3 – Pan et al. (2005) Table 2; 4 – Pan et al. (2005) Table 7; 5 – Hupe et al. (2012) Table 3; 6 – Sonnentrucker et al. (2007) Table 13; 7 – Velusamy et al. (2017) Table 2. aDensities derived from CN chemistry in Sheffer et al. (2008) and Pan et al. (2005) were recalculated using updated f -values and identical input parameter values, with the exception of HD 210121. See footnote f for details on HD 210121 input values. Table 5.2: —Continued
−3 a b Sight Line nH+H2 (T CO) (cm ) nH+H2 (CN) nH+H2 (C2) nH+H2 (C II) nH+H2 (C I) T H2 T C2 −3 −3 −3 −3 T CO(1-0) T CO(2-1) T CO(3-2) (cm ) (cm ) (cm ) (cm ) (K) (K) b Densities from C2 in Sonnentrucker et al. (2007) were recalculated using updated values for the cross sections that are larger by a factor of two, leading to a decrease of 50% in density. Densities
from C2 in Hupe et al. (2012) were recalculated using updated values for the cross sections that are larger by a factor of two and updated f -values, resulting in a total decrease of 60% in density. c T C2 values estimated in Sonnentrucker et al. (2007). d The density calculations from CN chemistry for HD 147888 includes N (C2) data from Sonnen- trucker et al. (2007) that was not included in the calculation of Pan et al. (2005). The recalculated −3 density from CN chemistry without the additional C2 data would have been 175 cm . e The density calculations from CN chemistry for HD 192035 includes N (C2) data from Hupe et al. 212 (2012) that was not included in the calculation of Sheffer et al. (2008). The recalculated density −3 from CN chemistry without the additional C2 data would have been 750 cm . f HD 210121 has an incorrect N (CH) value of 286.2 x 1012 cm−2 in Table 7 of Sheffer et al. (2008). The N (CH) value we used for HD 210121 is 28.61 x 1012 cm−2, which is consistent with the values listed for HD 210121 in Tables 3 and 4 of Sheffer et al. (2008). 5.2 GOT C+ Species Associations
Section 4.6.1 gave a component-by-component analysis for a single sight line and
Tables 4.13 and 4.14 showed component-by-component interpretations for individual
pointings. Section 4.6.2 provided similar results in summary. When the results of all
the pointings are combined, trends start to emerge. Table 5.3 shows the tabulated
results for all GOT C+ observations and we interpret the combined results here.
5.2.1 H I and Ca II Detections
The majority of components seen in H I are also seen in Ca II, which are both
probes of the more diffuse regions of individual clouds. Components only seen in Ca II
and H I gas are likely probing diffuse atomic gas, where densities are too low to excite
2 2 [C II] P3/2 → P1/2 efficiently or react rapidly enough to produce CO. This accounts
for approximately 1/4 of the associated components. Components seen in Ca II and
H I gas that are also associated with any of the additional species considered here
typically probe denser gas than the diffuse atomic regime.
5.2.2 C+, Ca I,K I, CH+, CH, and CN Detections
The majority of components seen in C+, Ca I,K I, CH+, CH, or CN also tend to be associated with probes of the very diffuse gas (H I and Ca II). The presence of C+,
Ca I,K I, CH+, or CH indicates that the component is probing denser material than if only Ca II or H I were observed (Pan et al. 2005). The additional presence of CN indicates that the component is probing even denser material than if only C+, Ca I,
K I, CH+, or CH were detected.
The species that probe very diffuse gas in emission seem to be associated with species observed in absorption that are seen in diffuse molecular clouds. The presence of C+, Ca I,K I, CH+, or CH all suggest that there should be some CO being produced,
213 as discussed in Section 4.6.1, and the presence of CN suggests even greater amounts
of CO should be present. Thus if there is no detectable amounts of CO emission at
the same Vlsr, it does not mean CO is not present, but that the conditions and/or density are not suitable for detecting the emission. Therefore, these components with
C+, Ca I,K I, CH+, or CH and no detectable CO are most likely probing regions of
CO dark gas, and components associated with CN and no detectable CO emission
are probably probing the slightly denser inner regions of CO dark gas.
These regions of CO dark gas can be further probed by looking at the pure ro-
tational emission lines of H2. These lines have been detected by the Infrared Space Observatory (ISO) and then by Spitzer several times – e.g. the Taurus molecular
cloud (Goldsmith et al. 2010a) and in translucent clouds (Ingalls et al. 2011a). The
James Webb Space Telescope (JWST) would likely build on these results using the
rotational emission lines of H2 to probe CO dark gas at the edges of molecular clouds.
5.2.3 CO Detections
Detectable amounts of CO emission often correspond to components seen in C+,
K I, CH+, CH, or CN. This accounts for almost half of the CO components. The other half of the associated CO emission components are only associated with Ca II compo- nents. About 70% of the CO emission components also have 13CO or C18O emission.
The additional presence of 13CO or C18O indicates a component that is associated with much denser gas, specifically the inner portion of the associated molecular cloud.
Detectable amounts of C18O are only present in the most dense regions considered here.
214 Table 5.3: Summary of GOT C+ Associations
Type of Gasa Diffuse Atomic Diffuse Molecular n∼100 cm−3 n∼300 cm−3 Total Additional Species Ca II K I or CH or CH+ CN and Associated Unassociated Detected only and Ca II Ca II or K I or CH or CH+ Components Emission Components Diffuse Atomic (WNM) H I only 7 12(5b) 1 20 31 (12b)
Diffuse Atomic (CNM) C+ only 1 1 0 2 5 (1c) C+ +H I 1 1 0 2 6 (1b)
Molecular gas 12CO only 1 0 1 2 3 C+ + 12CO 0 1(1c) 0 1 0 H I + 12CO 3 1 0 4 3 C+ + H I + 12CO 0 0 0 0 3 215 Molecular Cloud edge 12CO/13CO only 0 1 3 4 1 C+ + 12CO/13CO 2 0 1 3 0 H I + 12CO/13CO 2 1 0 3 3 C+ + H I + 12CO + 13CO 1 0 0 1 1
18CO + C+ or H I 2 0 0 2 3 or 12CO or 13CO
Total Associated 20 18 6 44 59 Components
% of Associated 45 % 41 % 14 % ...... Components
Additional Unassociated 37(8d) 9(3d) 0 46 . . . Absorption Components
a This table specifies the number of emission components that have been associated with at least one absorption component near the same Vlsr. b Number of included components without CO data. c Number of included components without H I data. d Number of included components that are also seen in a second nearby sight line. 5.2.4 Combined Absorption and Emission Picture
The majority of absorption components are only associated with gas seen in H I emission. About half of these components are representative of diffuse atomic gas, probed by Ca II, Ca I, and K I. The amount of C+ in atomic gas from the Warm
Neutral Medium is usually too low to produce C+ emission above the GOT C+ sensitivity limit. However, in the atomic and molecular gas of the Cold Neutral
Medium, C+ begins to increase above the GOT C+ sensitivity limit.
More importantly, the remaining half of the absorption components associated with H I only and/or C+ gas come from molecules found in diffuse molecular gas
+ + (e.g., Pan et al. 2005). In diffuse gas, reactions between C and H2 lead to CH and CH, and further reactions lead to CN (e.g., Federman et al. 1994) and CO (e.g.,
Sheffer et al. 2008), despite the amount of emission being below the sensitivity limit.
Therefore, components with C+, Ca I,K I, CH+, or CH and no CO emission are likely associated with CO dark gas.
There were very few detections of CN absorption in this sample, but when CN is seen, it is usually connected with gas containing CO emission. This further supports the combined picture summarized below. Furthermore, whenever these absorption components are associated with CO gas seen in emission, they are most likely probing the molecular cloud envelope. In the densest regions probed here, we also start to detect 13CO and C18O emission. In summary, we can say
• Diffuse Atomic Clouds have components seen in Ca II, Ca I,H I, and sometimes
+ −3 C ;(nH < 100 cm )
• Diffuse Molecular Clouds (CO dark gas) have components seen in C+, K I, CH+,
12 13 −3 −3 CH, or CN but without CO or CO emission; (100 cm ≤ nH ≤ 300 cm )
• Dense Molecular Cloud Envelopes have components seen in 12CO and 13CO
emission (with the distinction that clouds with 13CO are denser) and any of 216 −3 the additional species seen in emission or absorption. (300 cm ≤ nH ≤ 1000 cm−3)
5.3 CO emission and CO Dark Gas
The last two sections lay the path for the discussion of CO dark gas. All the discussed tracers of diffuse molecular clouds have been associated with detections of
12CO in absorption. Section 5.1 showed the densities inferred from the low J lines of 12CO absorption, C+ excitation, and CN chemistry all probe similar densities in diffuse clouds. Section 5.2 looks for kinematic associations between CO components in emission and the components in the species used in Section 5.1. Several of the species are observed throughout diffuse molecular clouds but are only associated with
CO emission for a handful of components. The lack of association of these tracers with CO in emission and the observed association with CO in absorption show that there is little if any difference in the characterization of diffuse molecular gas versus
CO dark gas. There are several well studied cloud complexes that have been observed over a range of wavelengths, and they help to demonstrate the implications for CO dark gas.
5.3.1 Chamaeleon Region
The Chamaeleon Region is one of the nearest well studied complexes of low mass star formation in the southern sky. The clouds in the Chamaeleon complex are at an approximate distance of 140 to 180 pc (Mizuno et al. 2001; Corradi et al. 2004).
Mizuno et al. (2001) found many filamentary structures throughout the region, with some of the more prominent filaments associated with dense (≥ 1000 cm−3) cloudlets found in the 13CO(J = 1-0) observations of Mizuno et al. (1998). Figure 5-1 shows CO emission contours in the Chamaeleon Region taken from the 12CO(J = 1-0) survey
217 with the NANTEN 4-meter radio telescope by Mizuno et al. (2001). The contour
level for the total integrated intensity of the 12CO(J = 1-0) line corresponds to a
3σ detection limit of 1.8 K km s−1. These contours have been augmented with data at UV and visible wavelengths from Andersson et al. (2002), Gry et al. (2002), and
Sheffer et al. (2007, 2008) as well as our own data toward HD 108927.
12 12 The column density of H2 and CO, as well as the kinetic and CO(J = 1-0) excitation temperatures, are given for each sight line. The isotopic ratio is also given for sight lines with 13CO detections. In all cases, CO absorption is detected from directions near or beyond the CO emission contours. The Vlsr of the dominant CO component is provided for each sight line. This can be compared with the average
Vlsr from CO emission, which is also shown in Figure 5-1. Comparing the velocities of the sight lines for the different complexes reveals likely associations.
Table 5.4 lists the V lsr of the observed component and the column densities of the
12 + H2, CO, CH , CH, and CN species for the sight lines in the region. The majority
−1 of the Chamaeleon sight lines have a component near the same V lsr of 3.4 km s , likely associated with the CO dark gas of the Cham II portion of the complex, seen
in emission. Similarly, the components at higher V lsr are part of the CO dark gas in
the Cham I portion of the complex, and components at lower V lsr are likely part of the CO dark gas in the Cham III portion of the complex. We suggest associations
with CO dark gas because the CO column densities derived from UV observations
are significantly lower than the values inferred from emission equivalent widths.
The largest H2 column density observed for these sight lines comes from the sight line toward HD 96675, just inside the CO emission contours of Chamaeleon I. This
sight line also had the largest gas density inferred from CN chemistry (400 cm−3).
However, the largest CO column density and gas density inferred from CO excitation
(400 cm−3) occurs toward HD 108927, which is just outside the emission contours of
Chamaeleon II and III. This sight line also has the largest CO/H2 column density
218 ratio, indicating a large amount of CO dark gas. Comparing this sight line to others,
we see 13CO detections near the emission contours, and as expected, we also observe
the amount of CO dark gas tends to decrease with distance from the emission contours.
The average gas density from CO excitation was around 140 cm−3 with two sight lines
(HD 116852 and HD 93237) indicating significantly lower values (60-70 cm−3).
The picture we introduced for the Chamaeleon Region complements other recent studies. Using C+ and HINSA features, Tang et al. (2016) found the fraction of dark molecular gas (DMG) in a cloud decreases with increasing excitation temperature and that DMG dominates cloud mass in diffuse gas. From observations of OH and
CO isotopologues, Xu et al. (2016) concluded that the dark molecular gas fraction
(DMGF) decreases from 80% in diffuse gas (AV = 0.4 mag) to 20% in denser gas
(AV = 2.6 mag) due to efficient CO formation. These trends were also observed in clouds near the Chamaeleon Region by Planck Collaboration et al. (2015) and Remy
et al. (2018b) using γ rays, and thermal dust emission along with H I and CO line
data. Remy et al. (2018b) also found that detectable CO emission in Chamaeleon
typically begins between a total gas column density of 0.6x1021 cm−2 and 2.5x1021
cm−2. Several of the sight lines from Table 5.4 have densities near or in this range.
219 iue51 h gr hw h oa nertditniycnor of contours intensity integrated total the shows figure The 5-1: Figure
0.5 pc at 150 pc r a enaatdfo iuoe l 20)adindicates and (2001) 3 al. the et to s equivalent Mizuno level from contour adapted a been has ure 12 h auspoie nTbe5.4. Table in provided values the − CO( 1 h aafralsgtlns xetH 097 oefrom come 108927, HD except lines, sight all for data The . J -)eiso nteCaalo ein h fig- The Region. Chamaeleon the in emission 1-0) = 220 σ eeto ii f18Kkm K 1.8 of limit detection Table 5.4: Chamaeleon Data
+ Sight Line Vlsr N H2 N CO N CH N CH N CN Reference (km s−1) (1019 cm−2) (1013 cm−2) (1012 cm−2) (1012 cm−2) (1012 cm−2) HD 93237 3.5 6.3(1.3) 2.5(0.5) ... 1.2(0.2) ... Sheffer et al. (2008) 3.6(0.5) ...... 1.2(1.2) ... Andersson et al. (2002) HD 94454 3.6 57.5(11.5) 20.0(4.0) ... 9.3(1.9) ... Sheffer et al. (2008) 3.7(0.1) ...... 9.4(1.0) ... Andersson et al. (2002) HD 96675 4.1 72.4(14.5) 190.5(38.1) 4.9(1.0) 22.9(4.6) 6.3(1.3) Sheffer et al. (2008) ... 70.0(15.0) ≥100 2.8 22.0 ... Gry et al. (2002) HD 99872 3.2 33.1(6.6) 44.7(8.9) 22.9(4.6) 12.9(2.6) ... Sheffer et al. (2008) 3.2(0.1) ...... 12.6(0.8) ... Andersson et al. (2002) HD 102065 1.0 ...... 5.0(1.0) 1.1(0.2) ... Sheffer et al. (2008) 221 3.8 36.3(7.3) 4.9(1.0) 5.8(1.2) 6.0(1.2) ... Sheffer et al. (2008) ... 34.0(3.0) 7.0 12.0 6.3 ... Gry et al. (2002) HD 108927 ... 32.0(3.0) ...... 14.0 ... Gry et al. (2002) 2.1(0.1) ...... 10.8(4.6) ... Andersson et al. (2002) HD 116852 0.6 6.8(1.4) 2.0(0.4) 1.9(0.3) 1.7(0.2) ... Sheffer et al. (2008) 5.3.2 Cepheus OB2 Region
The Cepheus OB2 Region is part of a ring shaped cloud system called the Cepheus bubble, at an approximate distance of 615 pc (de Zeeuw et al. 1999). Patel et al. (1998) proposed the bubble formed from stellar winds and supernovae explosions. Figure 5-2 shows CO emission contours of the Cepheus OB2 Region, adapted from Patel et al.
12 (1998) and Pan et al. (2005). The column densities of H2 and CO toward fourteen sight lines are provided. The data at UV wavelengths, shown in Figure 5-2, has been
12 taken from Pan et al. (2004, 2005) and Sheffer et al. (2007, 2008). N (H2) and N ( CO) are given in units of 1020 cm−2 and 1014 cm−2, respectively. N (12CO)/N (13CO) is provided when 13CO is detected. The densities from CN chemistry for these sight lines were given previously in Tables 4.9 and 4.10, and are primarily between 100 cm−3 and 250 cm−3. CO excitation temperatures and inferred densities were only available for a few of the sight lines in this region and are included in Table 5.2.
The inferred CO densities are similar to those from CN chemistry, consistent with our overall conclusions. 13CO was detected in sight lines with higher 12CO column densities, with the highest 12CO column densities found very near or inside the CO emission contours.
222 Figure 5-2: The figure shows 12CO emission contours of the Cep OB2 Re- gion, adapted from Patel et al. (1998); Pan et al. (2005). N(H2) and N (CO) are given for each sight line. N (12CO)/N (13CO) is 13 12 provided when CO is detected. N (H2) and N ( CO) have units of 1020 cm−2 and 1014 cm−2, respectively.
223 5.3.3 Perseus B5 Region
The Perseus cloud complex consists of an extended dusty H I envelope surrounding several regions of denser H2 gas (Sancisi et al. 1974; Sargent 1979; Ungerechts & Thaddeus 1987; Ridge et al. 2006). B5 is a relatively small cloud in the Perseus cloud complex, at an approximate distance of 260 pc (Cernis 1993). Figure 5-3 shows the 12CO 2.6 mm emission contours in the Perseus B5 Region. These contours were adapted from Wannier et al. (1999). The region observed in 2.6 mm emission is shown as a dashed line and several sight lines with visible and UV observations in or near this region are indicated. There are two sight lines slightly outside the region observed in emission, HD 24398 and HD 24534. The data in the figure come from
12 Sheffer et al. (2007) and Sheffer et al. (2008). N(H2) and N ( CO) are given for five sight lines in the region and N (12CO)/N (13CO) is provided when 13CO is detected.
Almost all of these sight lines lay outside the CO emission contours and are probing
−1 CO dark gas. The gas shown by the emission contours has a Vlsr around 10 km s . The gas sampled in each sight line has a very similar velocity component structure,
a primary component near 7 km s−1 and usually a secondary component near 5 km
s−1. This suggests that the common components in these sight lines are likely all
sampling different parts of the same structure. The densities from each of the sight
lines can be found in Table 5.2, revealing similar densities from CN chemistry and
CO(1-0) excitation, mainly between 125 and 320 cm−3.
There have been other proposed explanations for dark gas in the Perseus Molecular
Cloud. Andersson et al. (1992) mapped Perseus B5 in H I. Typically, 21 cm H I
emission is assumed to be optically thin. Fukui et al. (2014) and Fukui et al. (2015)
argue that dark gas can be explained by incorrect determinations of H I opacities and
find cold H I gas masses 2 to 2.5 times that of the optically thin calculations. However,
Li et al. (2015) analyzed the Perseus Molecular Cloud and found only a 10% increase
224 in the H I gas mass from optically thick H I calculations, not enough to account for the majority of the CO dark gas. Li et al. (2015) also find that the distributions of
CO dark gas and optically thick H I do not tend to follow each other. The search for cold optically thick H I in interstellar clouds by Reach et al. (2017b) suggests only a small amount of the dark gas can be accounted for by H I in cold atomic gas. They conclude that cold atomic gas is only a minor component of the total inferred dark gas. Recent work by Murray et al. (2018) confirm that optically thick H I does not dominate the dark gas in the local ISM and suggests the signatures of dark gas, seen as excess dust emission, are a combination of changes in dust emissivity and H2 in CO dark gas.
Optically thick H I has also been suggested by Reach et al. (2015) as one possible source of the variations seen in the gas to dust ratio. Further work by Reach et al.
(2017b) concludes that low values of the gas to dust ratio in diffuse clouds are the result of dust associated with H2 in CO dark gas, not optically thick H I. This is also consistent with an analysis of the break in slope of gas/reddening relationship
(N (H)/E(B − V )) by Liszt (2014). He finds that the break in slope at low E(B − V )
is due to the effects of H2 formation and cannot be attributed to the optical depth effects from optically thick H I.
225 Figure 5-3: The figure shows the 12CO 2.6 mm emission contours in the Perseus B5 Region Region, adapted from Wannier et al. (1999). The region observed in 2.6 mm emission is shown as a dashed 12 line. N(H2) and N ( CO) are given for each sight line in units of 1020 cm−2 and 1014 cm−2, respectively. N (12CO)/N (13CO) is provided when 13CO is detected.
226 5.3.4 Comparisons with Other Tracers
Several groups have found significant mass outside the core regions of CO emission, in translucent gas (Cotten et al. 2012) and in diffuse molecular gas (Liszt & Pety
2012). Cotten et al. (2012) found the OH 1667 MHz line to be an excellent tracer of gas in the low extinction regions of diffuse molecular gas. The equivalent width of the OH line (W (OH)) is linearly correlated with E(B − V ) in diffuse molecular clouds (Crutcher 1979; Grossman et al. 1990; Barriault et al. 2010; Cotten et al.
2012). Moreover, OH was determined to be closely connected with H2 and CH in translucent clouds from the UV/visible observations by Weselak et al. (2010) and in the reanalysis by Mookerjea (2016). Barriault et al. (2010) found that N (OH) increases monotonically with N (H I) up to N (OH) = 0.25 to 0.30 x 1014 cm−2 where
N (H I) saturates. This means N (H I) no longer increases monotonically at greater values of N (OH), further implying the presence of molecular gas not traced by H I
(Barriault et al. 2010).
HINSA features come from absorption from cold H I found mixed with molecu- lar gas in well shielded regions in front of a H I background emission source (Li &
Goldsmith 2003). Typical temperatures of the HINSA gas are less than 20 K (Gold- smith et al. 2007). H I narrow self absorption is referred to as “self” absorption to differentiate it from absorption against a background continuum source and “narrow” refers to the H I line width being narrower than the corresponding CO feature (Li &
Goldsmith 2003). OH components are seen at the same Vlsr as those in HINSA, C I, C18O, and 13CO and have similar line widths to HINSA and 13CO (Li & Goldsmith
2003; Goldsmith et al. 2007).
227 The results from tracers seen at UV and visible wavelengths such as CO, C I,C+,
CN, and C2 have already been discussed and can be found in Tables 5.1 and 5.2. Table 5.5 provides the information for additional species such as OH or OH+ when they are available. Clouds with CO emission contain measurable amounts of OH which falls off slower than the CO emission (Lewis et al. 1988), making OH an important probe of CO dark gas. Several authors have found a correspondence between OH observed at near UV wavelengths and both CH and CN (Weselak et al. 2009a; Porras et al. 2014), consistent with our results from CN chemistry and CO excitation. The sight lines with velocity information show correspondences between OH and CN and near ultraviolet OH+ absorption with CH+ (Porras et al. 2014). These measurements confirmed the results from submillimeter observations of OH+ in diffuse gas (Neufeld
et al. 2010).
There are several detections of OH and OH+ at UV/visible wavelengths in sight lines from the current data set as well as several of the companion data sets used in our analysis. Putting these observations into context, an average N(OH) ≈ 1.1 x 1014 was found in the survey of dark clouds by Li & Goldsmith (2003) and 0.5 x
1014 cm−2 in cirrus clouds by Barriault et al. (2010). Crutcher (1979) found a range
14 −2 −1 between 0.5 and 1 x 10 cm mag for N (OH)/AV in diffuse and dark clouds. This is in agreement with the value for the Chamaeleon I dark cloud of 1.5 x 1014 cm−2 mag−1 from Toriseva et al. (1985). Toriseva et al. (1985) made the first extensive radio line mapping of the Chamaeleon I dark cloud with 18 cm OH emission, finding
OH column densities that range from 15 to 76 x 1013 cm−2. Similar column densities
are found in diffuse molecular gas in the directions of background stars (see Table 5.5
for the reference cited there).
228 Several studies looked at OH in Perseus at radio (Sancisi et al. 1974; Wannier et al.
1993) and UV wavelengths (Federman et al. 1996; Roueff 1996; Felenbok & Roueff
−1 1996). These studies typically found OH at Vlsr between 8 and 11 km s , consistent with the structure seen in CN and CO absorption (Sheffer et al. 2008). In addition,
Wannier et al. (1993) included observations of strips near the Perseus B5 Region.
They found the OH emission to extend beyond the regions of CO emission, implying
the presence of OH throughout the diffuse molecular gas region, now referred to as
CO dark gas. The locations of the midpoint for the strips are provided in Table 5.5.
Rudnitskii (1978) first detected OH emission in Cepheus with components between
-7 and -16 km s−1, in agreement with the kinematics available from other species at the time (Wilson et al. 1974; Rydbeck et al. 1976). Moreover, Allen et al. (2012) conducted a small survey of OH emission near the Cepheus Region. It later became part of a larger pilot survey of dark molecular gas in the Galaxy focusing on the correspondence between H I, CO, and OH emission (Allen et al. 2015). They found
OH in 21/27 of their sight lines and no CO emission features without corresponding
OH emission and concluded OH to be a promising tracer for dark molecular gas.
Associating OH detections with the detections and non-detections of other species
can help further differentiate the various types of gas. Tang et al. (2017) combine
OH observations with C+ and H I data from GOT C+ and confirm that N (H I)
surrounding molecular clouds increases monotonically with N (OH). They categorize
gas with detections of OH and H I without CO emission as dark molecular gas (CO
dark gas). One of the next steps would be the incorporation of OH and OH+ into our combined picture of different tracers.
229 Table 5.5: Data for Additional Tracers
Sight Line Vlsr N OH N OH+ Reference (km s−1) (1013 cm−2) (1012 cm−2) Main Sample HD 149404 -19.2 ≤ 1.7 1.9 Porras et al. (2014) -12.8 ≤ 1.7 4.0 -7.0 ≤ 1.6 5.0 3.0 ≤ 2.5 3.0 1.1 4.3 ≤ 3.0 ...... 6.4 Krelowski et al. (2010) HD 151932 ... 7.9 ... Weselak et al. (2009a) ...... 3.2 Krelowski et al. (2010) HD 170740 ... 4.9 ... Weselak et al. (2009b) ... 5.4 ... Mookerjea (2016) ...... 2.3 Krelowski et al. (2010) Sheffer Sample HD 154368 -10.7 ≤ 2.0 1.7 Porras et al. (2014) -5.1 ≤ 2.4 4.2 -2.0 17.9 ≤ 2.9 0.5 ≤ 2.6 1.6 HD 210121 ... 7.1 2.4 Krelowski et al. (2010) Perseus Perseus B5 Average 8.4 10 ... Sancisi et al. (1974) HD 23180 ... 3.9 ... Roueff (1996) 10.8 ≤ 8.2 ... Federman et al. (1996) HD 24398 ... 4.1 ... Felenbok & Roueff (1996) 10.5 4.2 ... Federman et al. (1996) HD 24912 ... ≤ 5.5 ... Federman et al. (1996) Perseus Strips 3h28m42s 29◦2’20” 5.8 44 ... Wannier et al. (1993) 6.2 43 ... 5.5 29 ... 5.2 17 ... 3h16m37s 29◦47’ -0.1 45 ... Wannier et al. (1993) -0.2 53 ... 0.2 36 ... -0.1 21 ... 0.0 24 ... 0.6 23 ... 1.3 14 ... Chamaeleon ... 15.6-76.3 ... Toriseva et al. (1985)
230 5.4 Small Scale Structure toward h and χ Persei
The nearby double cluster, h and χ Persei, is among the most studied Galactic clusters in the northern hemisphere. Points et al. (2004) investigated the interstellar structure toward h and χ Persei in moderate and high resolution spectra of interstellar
Na I absorption and found evidence for large scale structure and individual sight line variations down to projected physical separations of 0.1 pc.
There are a number of sight lines in my sample toward h and χ Persei. The stars are at an average distance of approximately 2000 pc (see Table 2.1). Figure 5-4 shows the Ca II spectra for the h and χ Per sight lines. In the figure, the location of each sight line is shown in relation to an extinction contour (AV ≈ 0.5) of the region, indicating the concentration of denser material. The contours were modified from
figures provided in Dobashi (2011). Figures 5-5 and 5-6 are similar to Figure 5-4 but show CH and CH+ spectra.
M¨unch (1953, 1957) analyzed the velocity profiles from Na I and Ca II to search for evidence of large scale structure in interstellar clouds. He determined that the absorption spectra showed two distinct components, a grouping of faint narrow lines combining to form a V structure from the Perseus Arm at high blueshifted velocities and a saturated cluster of components associated with the Orion Spiral Arm observed at lower velocities. This was confirmed with higher resolution Na I observations by
Points et al. (2004).
All the sight lines shown in Figure 5-5 probe clouds in the Local Arm of the Milky
−1 Way. However, one component toward HD 13841 that has a V lsr at -42 km s and
−1 one component toward HD 16691 that has a V lsr at -47 km s , both observed in CH, probe the Perseus Arm, where the double cluster resides. The Perseus Arm is especially evident in Ca II, where there are several sight lines whose velocity compo- nents extend into this V lsr range. The structure from the Perseus Arm appears as a
231 second grouping of absorption features in Ca II. There also appears to be a potential
gradient in the amount of Ca II absorption for this second grouping as one moves across the cloud; there is no evidence of this structure from the Perseus Arm in the
Ca II spectrum toward directions in the Western part of the cluster (e.g., HD 19820).
The limited number of molecular components associated with the Perseus Arm likely is a result of prevailing low densities, as discussed above for the GOT C+ sample.
There are 5 members of h and χ Per included in the study by Sheffer et al. (2008).
They are HD 12323, HD 13268, HD 13745, HD 14434, and HD 15137. Three of the
sight lines have CN detections. The gas toward HD 13745 also has a component at
-44 km s−1 in CH, CH+, CN, and CO. That component is most likely probing the
same Perseus Arm seen in CH toward HD 13841 and HD 16691. Additionally, HD
13268 reveals an interstellar CH component at -35 km s−1. HD 13268 and HD 13745
were also observed by Ritchey et al. (2011). Their data illustrate Ca II structure
−1 extending to a V lsr of -80 km s and multiple groupings of interstellar absorption in
the spectra. In addition, Ritchey et al. (2011) present K I absorption at V lsr between -33 to -44 km s−1 for the two sight lines.
There is good agreement in velocity among nearby sight lines in our h and χ Per
sample. HD 13969, HD 14053, and BD+56◦ 0501 are all very close (≤ 0.5 pc at the
distance of h and χ Per) and show a common main interstellar component around
-9 km s−1 in CH absorption; this component is absent in the spectra toward HD
13841 despite the very close proximity. These four sight lines also have several minor
components in common in both CH and CH+. BD+56◦ 0563, BD+56◦ 0578, HD
14476, and HD 14443 are also all near each other (≤ 0.5 pc), revealing a common
component around -6 km s−1 in CH absorption. The sight lines toward HD 14443
and BD+56◦ 0563 also share a second component around -1 km s−1 that is absent in
the spectra toward BD+56◦ 0578 and HD 14476. Furthermore, there is a sight line
toward HD 14434 from Sheffer et al. (2008) very close (≤ 0.5 pc) to these directions
232 that has a CN component present at the same velocity (-1 km s−1). HD 17505 and HD
17520 are also near each other (≤ 0.5 pc), showing common components around -6
and -2 km s−1 in CH. However, in CH+ absorption there is agreement in the strength
of the -2 km s−1 component, but there is a much larger CH+ abundance toward HD
17505 in the -6 km s−1 component.
These differences among sight lines with such close proximity emphasize the im-
portance of small scale structure that extends well beyond the nearby dark cloud
structure shown as a contour of AV = 0.5. Almost all sight lines in this region seem
−1 −1 to have components at or near a V lsr of -12 km s and/or -2 km s in CH and CH+, indicating that the background stars in the double cluster are probing different portions of the same larger structure. This is especially true in Ca II, where the main components are usually near -12 km s−1 and -2 km s−1, but with larger uncertainties in V lsr associated with fitting complex and optically thick Ca II profiles. These results are consistent with the findings using high resolution Na I by Points et al. (2004); the individual absorption line components toward each sight line are likely not indi- vidual interstellar clouds but parts of a larger complex and that fluctuations in the structure and density of the gas could give rise to many of the variations observed in the interstellar absorption line profiles.
233 Figure 5-4: This figure shows the Ca II spectra for sight lines toward h and χ Per. The positions are overlaid on a map of extinction contours (AV = 0.5). The locations of the additional sight lines from Sheffer et al. (2008) are also indicated on the map.
234 Figure 5-5: This figure shows the CH spectra for sight lines toward h and χ Per. See the caption of Figure 5-4 for more details.
235 Figure 5-6: This figure shows the CH+ spectra for sight lines toward h and χ Per. See the caption of Figure 5-4 for more details.
236 Chapter 6
Conclusions
This work was intended to create a more unified picture of interstellar clouds
in the diffuse ISM. This was accomplished in part by bridging the gap between the
paradigm created by tracers seen in absorption at UV and visible wavelengths and
the one created by tracers seen in emission at longer wavelengths, allowing me to
describe the changing environment better than either could do alone.
In order to study cloud structure and physical conditions of the clouds, high
resolution observations at visible wavelengths of interstellar CN, CH, CH+, Ca I,
Ca II, and K I absorption were taken toward 52 stars, 35 in the primary sample and an additional 17 near GOT C+ pointings. After normalization, the absorption features in each spectrum were fit with a heavily modified version of the ISMOD code, a rms minimization fitting routine written in FORTRAN. The observations of sight lines in the primary sample were combined with archival UV spectra from FUSE and a few select supplementary observations from HST. The synthesized absorption profile of each component yielded information related to the physical conditions of the component (e.g., relative velocities, b-values, fractions and column densities for each J level). The additional sight lines with nearby GOT C+ targets were combined with C+ observations from Herschel, 12CO, 13CO, and C18O observations from the
ATNF Mopra Telescope, and H I 21 cm observations from the VLA Galactic Plane
237 Survey (Stil et al. 2006). The emission data for the GOT C+ sight lines were fit using
the MPFIT package from Markwardt (2009).
The species observed at visible wavelengths with high resolution spectra were fit
first. The majority of N (H2) is found in the low J levels (0 and 1). These low J levels of H2 are very optically thick, and most of the velocity component structure in those lines are blended together and cannot be reliably determined without addi- tional information. Since the archival UV data were observed at lower resolution, we used the well determined structure from lines of species at visible wavelengths. In particular, results from fits of CH absorption were adopted in synthesizing H2 data and results from CN, CH, and sometimes CH+ were the input for CO spectra. The resulting trends in the plot of CO versus H2 column density were used as a tool to explore changes in the chemistry and to make comparisons in density for the observed sight lines. Further insight came from plotting N (CH) against N (H2). The additional correspondences between components in emission and absorption helped to create a more unified picture of diffuse atomic and molecular gas in the interstellar medium.
The sight lines toward h and χ Per form a distinct subsample from the rest of the primary sample. In the h and χ Per sample, most of the CO is coming from low density CH+ like CH gas, while the Sheffer et al. (2008) sample has most of the gas coming from denser CN like CH gas. CO formation is less efficient in gas where CH+ is abundant. This results in less CH for the amount of CO present in the h and χ
Per sight lines, indicating a more diffuse local environment.
One of the objectives of this Dissertation was to determine the physical conditions in the environment between diffuse atomic gas and dense molecular gas. There are a wide range of densities between atomic and molecular clouds and only a complete set of density tracers can create a full picture. To this end, we utilized the column densities and excitation temperatures from CO and H2 to determine the gas density. The resulting gas densities from this method were compared to densities inferred
238 from other methods such as CN chemistry or C2 and C I excitation. These results were combined with those from kinematic correspondences between components in
emission and absorption. We found that species probing very diffuse gas in emission
seem to be associated with species observed in absorption that are seen in diffuse
molecular clouds. The picture described here is the latest effort seeking details of the
correspondences indicated by earlier studies.
A schematic of this combined picture is provided in Figure 6-1. This schematic
shows the inferred distributions of species across a plane parallel slab extending from
the cloud surface to its denser regions, highlighting the types of gas and species
observed in emission and absorption as red and black, respectively.
Diffuse atomic clouds, with densities less than ≈100 cm−3, typically have compo-
nents seen in Ca II, Ca I absorption, H I emission, and sometimes weak C+ emission.
In general, the gas densities derived from C I and C2 excitation are among the lowest densities probed here. These species occupy similar volumes of gas that have only small amounts of 12CO present.
Diffuse molecular clouds, with densities greater than ≈ 100 cm−3, have compo-
nents that are also seen in CH+, CH, K I, CN, and 12CO(J = 1-0) absorption as well
as C+ emission. From the CO chemistry, as discussed in Section 4.6.1, the presence
of C+, Ca I,K I, CH+, CH, or CN all indicate that there should be some CO being
produced, with the presence of CN indicative of even greater amounts of CO. The
majority of this CO is typically only seen in CO absorption from the low J levels
(i.e., CO dark gas). In other words, the majority of CO dark gas resides in diffuse
molecular clouds. These regions of CO dark gas can be further probed by looking
at the pure rotational emission lines of H2. These lines have been detected by ISO and then by Spitzer, and upcoming JWST observations would likely build on these results.
C+ is found throughout interstellar diffuse clouds. However, it is also sensitive
239 to pressure and the density of the gas and arises preferentially from denser regions.
On average, C+ emission from diffuse molecular gas resides between where the gas probed by CN chemistry and T CO(1-0) are located within the cloud. While there were only a few strong detections of CN absorption in my observations,
they were usually connected to gas containing CO emission. Typically, components
with strong C+ emission, strong CN absorption, excited lines of 12CO in absorption
and sometimes 12CO lines in emission are observed at higher gas densities, greater
than ≈ 300 cm−3. On average, the densities derived from CN chemistry fall between
those from T CO(1-0) and T CO(2-1), but are typically closer to the densities from
T CO(1-0). The CO gas being probed by the CO(1-0) transition and the CN molecule occupy similar volumes deeper into the diffuse molecular cloud. The densities from
T CO(2-1) indicate that the more excited CO resides deeper still. In the densest regions probed here, we also start to detect strong 12CO emission, as well as 13CO and C18O
emission. We plan to strengthen the picture described here through further analysis
of data for the GOT C+ project, including recently acquired spectra of CO and C I absorption with HST.
The OH radical has been found to be abundant in interstellar diffuse molecular gas and is a good tracer for dark molecular gas, coexisting with atomic and molecular hydrogen in regions without CO emission (Lewis et al. 1988; Magnani & Siskind 1990;
Wannier et al. 1993; Liszt & Lucas 1996; Goldsmith & Li 2005; Li et al. 2015; Xu et al. 2016; Mookerjea 2016). Correspondences between OH and other tracers already discussed (CH and CN) make it an ideal candidate to consider incorporating into our future work (Weselak et al. 2009a; Porras et al. 2014). Several groups have already combined observations of C+ and H I emission with OH to probe CO dark gas (Allen et al. 2012, 2015; Tang et al. 2017).
In addition, OH+ absorption features in the near ultraviolet are also correlated with CH+ features, providing a complimentary probe of more diffuse gas (Kre lowski
240 et al. 2010; Porras et al. 2014). The incorporation of OH and OH+ into our set of tracers would be a natural extension of this work and should reinforce the combined picture for CO dark gas.
The incorporation of triatomic species would also be an useful addition to this work. HCO+, in particular, has been detected in mm absorption in several studies and is a sensitive tracer of H2 gas (Liszt & Lucas 1994, 1996). In the recent work by Liszt et al. (2018), the 3 mm absorption lines of HCO+ were used to infer the content of the DNM in sight lines across the outer parts of Chamaeleon. They find that the
+ molecular H2 gas, inferred from HCO , can account for almost all the gas expected in previous DNM estimates for their sight lines from Planck Collaboration et al. (2015).
241 Figure 6-1: A schematic indicating the inferred distributions of species through the outer regions of an interstellar cloud, represented here by a plane parallel slab from the cloud surface to its denser regions. Different species are distributed according to gas den- sity. The red or black color indicates observations in emission or absorption, respectively. The spacing between hash-marks sym- bolizes changes in gas density (not to scale). Solid vertical lines indicate the boundaries between types of gas. Typical densities are given for the boundaries.
242 References
Abdo, A. A., Ackermann, M., Ajello, M., et al. 2010, ApJ, 710, 133
Ackermann, M., Ajello, M., Baldini, L., et al. 2011, ApJ, 726, 81
Aiello, S., Barsella, B., Chlewicki, G., et al. 1988, Astronomy and Astrophysics Sup-
plement Series, 73, 195
Allen, R. J., Hogg, D. E., & Engelke, P. D. 2015, AJ, 149, 123
Allen, R. J., Ivette Rodr´ıguez,M., Black, J. H., & Booth, R. S. 2012, AJ, 143, 97
Andersson, B. G., Roger, R. S., & Wannier, P. G. 1992, A&A, 260, 355
Andersson, B. G., Wannier, P. G., & Crawford, I. A. 2002, MNRAS, 334, 327
Bailer-Jones, C. A. L., Rybizki, J., Fouesneau, M., Mantelet, G., & Andrae, R. 2018,
AJ, 156, 58
Bally, J., & Langer, W. D. 1982, ApJ, 255, 143
Barbier, M. 1962, Journal des Observateurs, 45, 57
Barriault, L., Joncas, G., Lockman, F. J., & Martin, P. G. 2010, MNRAS, 407, 2645
Bhatt, N. H., & Cami, J. 2015, ApJS, 216, 22
Black, J. H., & Dalgarno, A. 1976, ApJ, 203, 132
—. 1977, The Astrophysical Journal Supplement Series, 34, 405
Black, J. H., & van Dishoeck, E. F. 1988, ApJ, 331, 986 243 Blitz, L., Bazell, D., & Desert, F. X. 1990, ApJ, 352, L13
Bolatto, A. D., Wolfire, M., & Leroy, A. K. 2013, Annual Review of Astronomy and
Astrophysics, 51, 207
Burgh, E. B., France, K., & McCandliss, S. R. 2007, ApJ, 658, 446
Buss, Richard H., J., Allen, M., McCandliss, S., et al. 1994, ApJ, 430, 630
Cannon, A. J., & Pickering, E. C. 1993, VizieR Online Data Catalog, 3135
Carruthers, G. R. 1970, ApJ, 161, L81
Casassus, S., Stahl, O., & Wilson, T. L. 2005, A&A, 441, 181
Cernis, K. 1993, Baltic Astronomy, 2, 214
Chan, W. F., Cooper, G., & Brion, C. E. 1993, Chemical Physics, 170, 123
Claria, J. J. 1974, AJ, 79, 1022
Cloutier, R., Currie, T., Rieke, G. H., et al. 2014, ApJ, 796, 127
Code, A. D., Davis, J., Bless, R. C., & Hanbury Brown, R. 1976, ApJ, 203, 417
Corradi, W. J. B., Franco, G. A. P., & Knude, J. 2004, MNRAS, 347, 1065
Cotten, D. L., & Magnani, L. 2013, MNRAS, 436, 1152
Cotten, D. L., Magnani, L., Wennerstrom, E. A., Douglas, K. A., & Onello, J. S.
2012, AJ, 144, 163
Cox, N. L. J., Ehrenfreund, P., Foing, B. H., et al. 2011, A&A, 531, A25
Crawford, I. A., Barlow, M. J., Diego, F., & Spyromilio, J. 1994, MNRAS, 266, 903
Crenny, T., & Federman, S. R. 2004, ApJ, 605, 278
244 Crutcher, R. M. 1979, ApJ, 234, 881
—. 1985, ApJ, 288, 604
Crutcher, R. M., & Federman, S. R. 1987, ApJ, 316, L71
Currie, T., Kenyon, S. J., Balog, Z., et al. 2008, ApJ, 672, 558
Currie, T., Hernandez, J., Irwin, J., et al. 2010, The Astrophysical Journal Supple-
ment Series, 186, 191
Danks, A. C., Federman, S. R., & Lambert, D. L. 1984, A&A, 130, 62
De Becker, M. 2013, Bulletin de la Societe Royale des Sciences de Liege, 82, 33 de Graauw, T., Helmich, F. P., Phillips, T. G., et al. 2010, A&A, 518, L6 de Zeeuw, P. T., Hoogerwerf, R., de Bruijne, J. H. J., Brown, A. G. A., & Blaauw,
A. 1999, AJ, 117, 354
Dekker, H., Delabre, B., D’Odorico, S., et al. 1986, The Messenger, 43, 27
Dickey, J. M., McClure-Griffiths, N., Gibson, S. J., et al. 2013, Publications of the
Astronomical Society of Australia, 30, e003
Dixon, W. V., Sahnow, D. J., Barrett, P. E., et al. 2007, Publications of the Astro-
nomical Society of the Pacific, 119, 527
Dobashi, K. 2011, PASJ, 63, S1
Dobashi, K., Bernard, J.-P., Yonekura, Y., & Fukui, Y. 1994, ApJS, 95, 419
Dobashi, K., Uehara, H., Kandori, R., et al. 2005, PASJ, 57, S1
Douglas, K. A., & Taylor, A. R. 2007, ApJ, 659, 426
245 Draine, B. T. 2011, Physics of the Interstellar and Intergalactic Medium (Princeton
University Press)
Ducati, J. R. 2002, VizieR Online Data Catalog, 2237
Dunham, M. K., Rosolowsky, E., Evans, II, N. J., Cyganowski, C., & Urquhart, J. S.
2011, ApJ, 741, 110
Erman, P., & Iwamae, A. 1995, ApJ, 450, L31
ESA, ed. 1997, ESA Special Publication, Vol. 1200, The HIPPARCOS and TYCHO
catalogues. Astrometric and photometric star catalogues derived from the ESA
HIPPARCOS Space Astrometry Mission
Federer, W., Villinger, H., Lindinger, W., & Ferguson, E. E. 1986, Chemical Physics
Letters, 123, 12
Federman, S. R. 1982, ApJ, 253, 601
Federman, S. R., Danks, A. C., & Lambert, D. L. 1984, ApJ, 287, 219
Federman, S. R., Fritts, M., Cheng, S., et al. 2001, ApJS, 134, 133
Federman, S. R., Glassgold, A. E., Jenkins, E. B., & Shaya, E. J. 1980, ApJ, 242, 545
Federman, S. R., & Huntress, W. T., J. 1989, ApJ, 338, 140
Federman, S. R., Huntress, W. T., J., & Prasad, S. S. 1990, ApJ, 354, 504
Federman, S. R., & Lambert, D. L. 1988, ApJ, 328, 777
Federman, S. R., Sheffer, Y., Lambert, D. L., & Smith, V. V. 2005, ApJ, 619, 884
Federman, S. R., Strom, C. J., Lambert, D. L., et al. 1994, ApJ, 424, 772
Federman, S. R., Weber, J., & Lambert, D. L. 1996, ApJ, 463, 181
246 Federman, S. R., & Willson, R. F. 1982, ApJ, 260, 124
—. 1984, ApJ, 283, 626
Felenbok, P., & Roueff, E. 1996, ApJ, 465, L57
Fitzpatrick, E. L., & Massa, D. 2007, ApJ, 663, 320
Franco, M. L., Magazzu, A., & Stalio, R. 1985, A&A, 147, 191
Fukui, Y., Torii, K., Onishi, T., et al. 2015, ApJ, 798, 6
Fukui, Y., Okamoto, R., Kaji, R., et al. 2014, ApJ, 796, 59
Gaia Collaboration, Clementini, G., Eyer, L., et al. 2017, A&A, 605, A79
Garmany, C. D., & Stencel, R. E. 1992, Astronomy and Astrophysics Supplement
Series, 94, 211
Geminale, A., & Popowski, P. 2004, Acta Astron., 54, 375
Glassgold, A. E., & Langer, W. D. 1974, ApJ, 193, 73
—. 1976, ApJ, 206, 85
Goldsmith, P. F. 2013, ApJ, 774, 134
Goldsmith, P. F., & Li, D. 2005, ApJ, 622, 938
Goldsmith, P. F., Li, D., & Krˇco,M. 2007, ApJ, 654, 273
Goldsmith, P. F., Velusamy, T., Li, D., & Langer, W. D. 2010a, ApJ, 715, 1370
—. 2010b, ApJ, 715, 1370
Gordon, K. D., Cartledge, S., & Clayton, G. C. 2009, ApJ, 705, 1320
Gredel, R., van Dishoeck, E. F., & Black, J. H. 1991, A&A, 251, 625 247 —. 1993, A&A, 269, 477
—. 1994, A&A, 285, 300
Gredel, R., van Dishoeck, E. F., de Vries, C. P., & Black, J. H. 1992, A&A, 257, 245
Green, S., & Thaddeus, P. 1976, ApJ, 205, 766
Grenier, I. A., Casandjian, J.-M., & Terrier, R. 2005, Science, 307, 1292
Grossman, V., Heithausen, A., Meyerdierks, H., & Mebold, U. 1990, A&A, 240, 400
Gry, C., Boulanger, F., Nehm´e,C., et al. 2002, A&A, 391, 675
Guetter, H. H. 1968, PASP, 80, 197
Hamann, W. R., Gr¨afener,G., & Liermann, A. 2006, A&A, 457, 1015
Hanson, M. M., & Clayton, G. C. 1993, AJ, 106, 1947
Hanson, M. M., Clayton, G. C., Anderson, C. M., & Veen, K. A. 1994, in The First
Symposium on the Infrared Cirrus and Diffuse Interstellar Clouds. ASP Conference
Series, Vol. 58, 1994, R.M. Cutri and W.B. Latter, Eds., p.84, Vol. 58, 84
Herbst, E. 1978, ApJ, 222, 508
Hiltner, W. A., Garrison, R. F., & Schild, R. E. 1969, ApJ, 157, 313
Hirschauer, A., Federman, S. R., Wallerstein, G., & Means, T. 2009, ApJ, 696, 1533
Høg, E., Fabricius, C., Makarov, V. V., et al. 2000a, A&A, 355, L27
—. 2000b, A&A, 355, L27
Hohle, M. M., Neuh¨auser,R., & Schutz, B. F. 2010, Astronomische Nachrichten, 331,
349
248 Hollenbach, D. J., Werner, M. W., & Salpeter, E. E. 1971, ApJ, 163, 165
Houk, N. 1978, Michigan catalogue of two-dimensional spectral types for the HD stars
(University Microfilms International)
Houk, N., & Cowley, A. P. 1975, University of Michigan Catalogue of two-dimensional
spectral types for the HD stars. Volume I. Declinations -90 to -53 (University Mi-
crofilms International)
Houk, N., & Smith-Moore, M. 1988, Michigan Catalogue of Two-dimensional Spectral
Types for the HD Stars. Declinations -26.0 deg to -12.0 deg., Vol. 4 (Department
of Astronomy, University of Michigan, Ann Arbor)
Houk, N., & Swift, C. 1999, Michigan Spectral Survey, 5, 0
Humphreys, R. M. 1978, The Astrophysical Journal Supplement Series, 38, 309
Hupe, R. C., Sheffer, Y., & Federman, S. R. 2012, ApJ, 761, 38
Ingalls, J. G., Bania, T. M., Boulanger, F., et al. 2011a, ApJ, 743, 174
—. 2011b, ApJ, 743, 174
Ingalls, J. G., Chamberlin, R. A., Bania, T. M., et al. 1997, ApJ, 479, 296
Jaschek, M., & Egret, D. 1982, in IAU Symposium, Vol. 98, Be Stars, ed. M. Jaschek
& H.-G. Groth, 261
Jenkins, E. B., & Tripp, T. M. 2011, ApJ, 734, 65
Jenniskens, P., & Greenberg, J. M. 1993, A&A, 274, 439
Jensen, A. G., Snow, T. P., Sonneborn, G., & Rachford, B. L. 2010, ApJ, 711, 1236
Kavars, D. W., Dickey, J. M., McClure-Griffiths, N. M., Gaensler, B. M., & Green,
A. J. 2005, ApJ, 626, 887 249 Ka´zmierczak,M., Schmidt, M. R., Bondar, A., & Kre lowski, J. 2010, MNRAS, 402,
2548
Kim, B. G., Kawamura, A., Yonekura, Y., & Fukui, Y. 2004, PASJ, 56, 313
Kislyuk, V., Yatsenko, A., Ivanov, G., Pakulyak, L., & Sergeeva, T. 2000, in Journ´ees
1999 - syst`emesde r´ef´erencespatio-temporels IX. Lohrmann-Kolloquium. Motion
of celestial bodies, astrometry and astronomical reference frames, ed. M. Soffel &
N. Capitaine, 61
Knapp, G. R., & Jura, M. 1976, ApJ, 209, 782
Knauth, D. C. 2001, PhD thesis, The University of Toledo
Kos, J., & Zwitter, T. 2013, ApJ, 774, 72
Kre lowski, J., Beletsky, Y., & Galazutdinov, G. A. 2010, ApJ, 719, L20
Kre lowski, J., Ehrenfreund, P., Foing, B. H., et al. 1999, A&A, 347, 235
Kre lowski, J., Galazutdinov, G., & Gnaci´nski,P. 2012, Astronomische Nachrichten,
333, 627
Krelowski, J., & Strobel, A. 1983, A&A, 127, 271
Lambert, D. L., Sheffer, Y., & Crane, P. 1990, ApJ, 359, L19
Lambert, D. L., Sheffer, Y., Gilliland, R. L., & Federman, S. R. 1994, ApJ, 420, 756
Lang, K. R., & Willson, R. F. 1978, ApJ, 224, 125
Langer, W. D., Velusamy, T., Pineda, J. L., et al. 2010, A&A, 521, L17
Langer, W. D., Velusamy, T., Pineda, J. L., Willacy, K., & Goldsmith, P. F. 2014,
A&A, 561, A122
250 Langhoff, S. R., Bauschlicher, Charles W., J., Rendell, A. P., & Komornicki, A. 1990,
J. Chem. Phys., 92, 6599
Le Petit, F., Nehm´e,C., Le Bourlot, J., & Roueff, E. 2006, The Astrophysical Journal
Supplement Series, 164, 506
Lennon, D. J., Dufton, P. L., & Fitzsimmons, A. 1992, Astronomy and Astrophysics
Supplement Series, 94, 569
Lewis, B. M., Federman, S. R., & Shaya, E. 1988, in Bulletin of the American Astro-
nomical Society, Vol. 20, 642
Li, D., & Goldsmith, P. F. 2003, ApJ, 585, 823
Li, D., Xu, D., Heiles, C., Pan, Z., & Tang, N. 2015, Publication of Korean Astro-
nomical Society, 30, 75
Li, D., Tang, N., Nguyen, H., et al. 2018, ApJS, 235, 1
Liszt, H. 2014, ApJ, 783, 17
Liszt, H., Gerin, M., & Grenier, I. 2018, ArXiv e-prints, arXiv:1807.02210
Liszt, H., & Lucas, R. 1996, A&A, 314, 917
Liszt, H. S. 1979, ApJ, 233, L147
—. 2011, A&A, 527, A45
Liszt, H. S., & Lucas, R. 1994, ApJ, 431, L131
—. 1998, A&A, 339, 561
Liszt, H. S., & Pety, J. 2012, A&A, 541, A58
Liszt, H. S., Pety, J., & Lucas, R. 2010, A&A, 518, A45
251 Lutz, T. E., & Lutz, J. H. 1977, AJ, 82, 431
Lynds, B. T. 1962, ApJS, 7, 1
Magnani, L., Chastain, R. J., Kim, H. C., et al. 2003, ApJ, 586, 1111
Magnani, L., Lada, E. A., Sandell, G., & Blitz, L. 1989, ApJ, 339, 244
Magnani, L., & Onello, J. S. 1993, ApJ, 408, 559
—. 1995, ApJ, 443, 169
Magnani, L., Onello, J. S., Adams, N. G., Hartmann, D., & Thaddeus, P. 1998, ApJ,
504, 290
Magnani, L., & Siskind, L. 1990, ApJ, 359, 355
Markwardt, C. B. 2009, in Astronomical Society of the Pacific Conference Series, Vol.
411, Astronomical Data Analysis Software and Systems XVIII, ed. D. A. Bohlender,
D. Durand, & P. Dowler, 251
Marquette, J. B., Rowe, B. R., Dupeyrat, G., Poissant, G., & Rebrion, C. 1985,
Chemical Physics Letters, 122, 431
Martins, F., Escolano, C., Wade, G. A., et al. 2012, A&A, 538, A29
Massa, D., & Savage, B. D. 1984, ApJ, 279, 310
Mathys, G. 1989, Astronomy and Astrophysics Supplement Series, 81, 237
Mattila, K. 1986, A&A, 160, 157
McElroy, D., Walsh, C., Markwick, A. J., et al. 2013, A&A, 550, doi:10.1051/0004-
6361/201220465
Megier, A., Strobel, A., Bondar, A., et al. 2005, ApJ, 634, 451
252 Megier, A., Strobel, A., Galazutdinov, G. A., & Kre lowski, J. 2009, A&A, 507, 833
Mizuno, A., Yamaguchi, R., Tachihara, K., et al. 2001, Publications of the Astro-
nomical Society of Japan, 53, 1071
Mizuno, A., Hayakawa, T., Yamaguchi, N., et al. 1998, ApJ, 507, L83
Mookerjea, B. 2016, MNRAS, 459, 2822
Morgan, W. W., Code, A. D., & Whitford, A. E. 1955, ApJS, 2, 41
M¨unch, G. 1953, Publications of the Astronomical Society of the Pacific, 65, 179
—. 1957, ApJ, 125, 42
Murray, C. E., Peek, J. E. G., Lee, M.-Y., & Stanimirovic, S. 2018, ArXiv e-prints,
arXiv:1806.01300
National Optical Astronomy Observatories. 1999, IRAF: Image Reduction and Anal-
ysis Facility, Astrophysics Source Code Library, ascl:9911.002
Neufeld, D. A., Goicoechea, J. R., Sonnentrucker, P., et al. 2010, A&A, 521, L10
Palazzi, E., Mandolesi, N., & Crane, P. 1992, ApJ, 398, 53
Pan, K. 2002, PhD thesis, THE UNIVERSITY OF TOLEDO
Pan, K., Federman, S. R., Cunha, K., Smith, V. V., & Welty, D. E. 2004, ApJS, 151,
313
Pan, K., Federman, S. R., Sheffer, Y., & Andersson, B.-G. 2005, ApJ, 633, 986
Pan, K., Federman, S. R., & Welty, D. E. 2001, ApJ, 558, L105
Papaj, J., & Krelowski, J. 1992, Acta Astron., 42, 211
Papaj, J., Wegner, W., & Krelowski, J. 1991, MNRAS, 252, 403 253 Paradis, D., Dobashi, K., Shimoikura, T., et al. 2012, A&A, 543, A103
Patel, N. A., Goldsmith, P. F., Heyer, M. H., Snell, R. L., & Pratap, P. 1998, ApJ,
507, 241
Patriarchi, P., Morbidelli, L., & Perinotto, M. 2003, A&A, 410, 905
Patriarchi, P., & Perinotto, M. 1999, Memorie della Societa Astronomica Italiana,
70, 637
Pilbratt, G. L., Riedinger, J. R., Passvogel, T., et al. 2010, A&A, 518, L1
Pineda, J. L., Langer, W. D., Velusamy, T., & Goldsmith, P. F. 2013, A&A, 554,
A103
Pineda, J. L., Velusamy, T., Langer, W. D., et al. 2010, A&A, 521, L19
Planck Collaboration, Fermi Collaboration, Ade, P. A. R., et al. 2015, A&A, 582,
A31
Points, S. D., Lauroesch, J. T., & Meyer, D. M. 2004, Publications of the Astronomical
Society of the Pacific, 116, 801
Porras, A. J., Federman, S. R., Welty, D. E., & Ritchey, A. M. 2014, ApJ, 781, L8
Rachford, B. L., Snow, T. P., Tumlinson, J., et al. 2002, ApJ, 577, 221
Rachford, B. L., Snow, T. P., Destree, J. D., et al. 2009, The Astrophysical Journal
Supplement Series, 180, 125
Reach, W. T., Bernard, J.-P., Jarrett, T. H., & Heiles, C. 2017a, ApJ, 851, 119
Reach, W. T., Heiles, C., & Bernard, J.-P. 2015, ApJ, 811, 118
—. 2017b, ApJ, 834, 63
254 Reach, W. T., Koo, B.-C., & Heiles, C. 1994, ApJ, 429, 672
Reach, W. T., Wall, W. F., & Odegard, N. 1998, ApJ, 507, 507
Rebrion, C., Marquette, J. B., Rowe, B. R., & Clary, D. C. 1988, Chemical Physics
Letters, 143, 130
Reed, B. C. 2003, AJ, 125, 2531
Remy, Q., Grenier, I. A., Marshall, D. J., & Casandjian, J. M. 2017, A&A, 601, A78
—. 2018a, ArXiv e-prints, arXiv:1803.05686
—. 2018b, A&A, 611, A51
Rice, J. S., Federman, S. R., Flagey, N., et al. 2018, ApJ, 858, 111
Ridge, N. A., Schnee, S. L., Goodman, A. A., & Foster, J. B. 2006, ApJ, 643, 932
Ritchey, A. M. 2009, PhD thesis, The University of Toledo
Ritchey, A. M., Federman, S. R., & Lambert, D. L. 2011, ApJ, 728, 36
Roueff, E. 1996, MNRAS, 279, L37
Rudnitskii, G. M. 1978, AZh, 55, 345
Rydbeck, O. E. H., Kollberg, E., Hjalmarson, A.,˚ et al. 1976, The Astrophysical
Journal Supplement Series, 31, 333
Saesen, S., Carrier, F., Pigulski, A., et al. 2010, A&A, 515, A16
Sancisi, R., Goss, W. M., Andersson, C., Johansson, B., & Winnberg, A. 1974, A&A,
35, 445
Sargent, A. I. 1979, ApJ, 233, 163
255 Savage, B. D., Bohlin, R. C., Drake, J. F., & Budich, W. 1977, ApJ, 216, 291
Savage, B. D., Massa, D., Meade, M., & Wesselius, P. R. 1985, The Astrophysical
Journal Supplement Series, 59, 397
Schlingman, W. M., Shirley, Y. L., Schenk, D. E., et al. 2011, ApJS, 195, 14
Sheffer, Y., Federman, S. R., & Andersson, B.-G. 2003, ApJ, 597, L29
Sheffer, Y., Rogers, M., Federman, S. R., et al. 2008, ApJ, 687, 1075
Sheffer, Y., Rogers, M., Federman, S. R., Lambert, D. L., & Gredel, R. 2007, ApJ,
667, 1002
Shenavrin, V. I., Taranova, O. G., & Nadzhip, A. E. 2011, Astronomy Reports, 55,
31
Shepler, B. C., Yang, B. H., Dhilip Kumar, T. J., et al. 2007, A&A, 475, L15
S lyk, K., Bondar, A. V., Galazutdinov, G. A., & Kre lowski, J. 2008, MNRAS, 390,
1733
Smith, A. M., & Stecher, T. P. 1971, ApJ, 164, L43
Snow, T. P., & McCall, B. J. 2006, Annual Review of Astronomy and Astrophysics,
44, 367
Sofia, U. J., Cardelli, J. A., Guerin, K. P., & Meyer, D. M. 1997, ApJ, 482, L105
Sonnentrucker, P., Welty, D. E., Thorburn, J. A., & York, D. G. 2007, The Astro-
physical Journal Supplement Series, 168, 58
Sota, A., Ma´ızApell´aniz, J., Morrell, N. I., et al. 2014, The Astrophysical Journal
Supplement Series, 211, 10
256 Sota, A., Ma´ızApell´aniz,J., Walborn, N. R., et al. 2011, The Astrophysical Journal
Supplement Series, 193, 24
Spitzer, L., J., & Jenkins, E. B. 1975, Annual Review of Astronomy and Astrophysics,
13, 133
Spitzer, Lyman, J., & Zabriskie, F. R. 1959, Publications of the Astronomical Society
of the Pacific, 71, 412
Stahl, O., Casassus, S., & Wilson, T. 2008, A&A, 477, 865
Stil, J. M., Taylor, A. R., Dickey, J. M., et al. 2006, AJ, 132, 1158
Tang, N., Li, D., Heiles, C., et al. 2016, A&A, 593, A42
—. 2017, ArXiv e-prints, arXiv:1703.05864
The, P. S., de Winter, D., Arens, M., Heijblok, M., & Nieuwland, E. R. 1989, As-
tronomy and Astrophysics Supplement Series, 81, 115
Thorburn, J. A., Hobbs, L. M., McCall, B. J., et al. 2003, ApJ, 584, 339
Tielens, A. G. G. M. 2005, The Physics and Chemistry of the Interstellar Medium
(Princeton University Press)
Togi, A., & Smith, J. D. T. 2016, ApJ, 830, 18
Toriseva, M., H¨oglund,B., & Mattila, K. 1985, Revista Mexicana de Astronomia y
Astrofisica, vol. 10, 10, 135
Tull, R. G. 1998, in Optical Astronomical Instrumentation, Vol. 3355, 387–398
Tull, R. G., MacQueen, P. J., Sneden, C., & Lambert, D. L. 1995, PASP, 107, 251
Ungerechts, H., & Thaddeus, P. 1987, The Astrophysical Journal Supplement Series,
63, 645 257 Vacca, W. D., & Torres-Dodgen, A. V. 1990, The Astrophysical Journal Supplement
Series, 73, 685 van Breda, I. G., & Whittet, D. C. B. 1977, Ap&SS, 48, 297 van Dishoeck, E. F., & Black, J. H. 1988, ApJ, 334, 771 van Dishoeck, E. F., & de Zeeuw, T. 1984, MNRAS, 206, 383 van Leeuwen, F. 2007, A&A, 474, 653
Vejby-Christensen, L., Andersen, L. H., Heber, O., et al. 1997, ApJ, 483, 531
Velusamy, T., Langer, W. D., Goldsmith, P. F., & Pineda, J. L. 2017, ApJ, 838, 165
Velusamy, T., Langer, W. D., Pineda, J. L., et al. 2010, A&A, 521, L18
Viggiano, A. A., Howorka, F., Albritton, D. L., & Fehsenfeld, F. C. 1980, ApJ, 236,
492
Walborn, N. R. 1971, The Astrophysical Journal Supplement Series, 23, 257
Walborn, N. R., & Fitzpatrick, E. L. 2000, PASP, 112, 50
Wannier, P., Andersson, B. G., Penprase, B. E., & Federman, S. R. 1999, ApJ, 510,
291
Wannier, P. G., Andersson, B. G., Federman, S. R., et al. 1993, ApJ, 407, 163
Watson, W. D., Anicich, V. G., & Huntress, Jr., W. T. 1976, ApJ, 205, L165
Weaver, H., Williams, D. R. W., Dieter, N. H., & Lum, W. T. 1965, Nature, 208, 29
Wegner, W. 2002, Baltic Astronomy, 11, 1
—. 2003, Astronomische Nachrichten, 324, 219
258 Weinreb, S., Barrett, A. H., Meeks, M. L., & Henry, J. C. 1963, Nature, 200, 829
Welsh, B. Y., & Sallmen, S. 2003, A&A, 408, 545
Welty, D. E., Federman, S. R., Gredel, R., Thorburn, J. A., & Lambert, D. L. 2006,
The Astrophysical Journal Supplement Series, 165, 138
Welty, D. E., & Hobbs, L. M. 2001, ApJS, 133, 345
Welty, D. E., Hobbs, L. M., & Morton, D. C. 2003, ApJS, 147, 61
Welty, D. E., Morton, D. C., & Hobbs, L. M. 1996, ApJS, 106, 533
Wenger, M., Ochsenbein, F., Egret, D., et al. 2000, A&AS, 143, 9
Weselak, T., Galazutdinov, G., Beletsky, Y., & Kre lowski, J. 2009a, A&A, 499, 783
Weselak, T., Galazutdinov, G. A., Beletsky, Y., & Kre lowski, J. 2009b, MNRAS, 400,
392
—. 2010, MNRAS, 402, 1991
Weselak, T., Galazutdinov, G. A., Musaev, F. A., Beletsky, Y., & Kre lowski, J. 2009c,
A&A, 495, 189
Weselak, T., Galazutdinov, G. A., Musaev, F. A., & Kre lowski, J. 2008, A&A, 484,
381
Weselak, T., Galazutdinov, G. A., Sergeev, O., et al. 2014, Acta Astron., 64, 371
Weselak, T., & Kre lowski, J. 2014, ArXiv e-prints, arXiv:1408.3071
Whittet, D. C. B., van Breda, I. G., & Glass, I. S. 1976, MNRAS, 177, 625
Williams, P. M., Kidger, M. R., van der Hucht, K. A., et al. 2001, MNRAS, 324, 156
Willson, R. F. 1981, ApJ, 247, 116 259 Wilson, R. W., Jefferts, K. B., & Penzias, A. A. 1970, ApJ, 161, L43
Wilson, T. L., Mauersberger, R., Gensheimer, P. D., Muders, D., & Bieging, J. H.
1999, ApJ, 525, 343
Wilson, W. J., Schwartz, P. R., Epstein, E. E., et al. 1974, ApJ, 191, 357
Wolfire, M. G., Hollenbach, D., & McKee, C. F. 2010, ApJ, 716, 1191
Wouterloot, J. G. A. 1981, PhD thesis, -
Wouterloot, J. G. A., & Brand, J. 1989, A&AS, 80, 149
Wouterloot, J. G. A., Walmsley, C. M., & Henkel, C. 1988, A&A, 203, 367
Xu, D., & Li, D. 2016, ApJ, 833, 90
Xu, D., Li, D., Yue, N., & Goldsmith, P. F. 2016, ApJ, 819, 22
Yang, B., Stancil, P. C., Balakrishnan, N., Forrey, R. C., & Bowman, J. M. 2013,
ApJ, 771, 49
Zacharias, N., Finch, C. T., Girard, T. M., et al. 2012, VizieR Online Data Catalog,
I/322A
Zdanaviˇcius,J., & Zdanaviˇcius,K. 2002, Baltic Astronomy, 11, 441
Zhang, X.-B., Zhang, R.-X., & Li, Z.-P. 2005, Chinese J. Astron. Astrophys., 29, 9Z
Zsarg´o,J., & Federman, S. R. 2003, ApJ, 589, 319
260 Appendix A
Spectra from Ground-based
Observations
The spectra shown here come from McDonald observations and are described in
Section 3.2. See Figures 3-1 and 3-2 for additional details about the plots of these spectra.
261 Figure A-1: Stacked spectra of McDonald data toward HD 108.
262 Figure A-2: Stacked spectra of McDonald data toward HD 5689.
263 Figure A-3: Stacked spectra of McDonald data toward HD 12882.
264 Figure A-4: Stacked spectra of McDonald data toward HD 13969.
265 Figure A-5: Stacked spectra of McDonald data toward HD 14053.
266 Figure A-6: Stacked spectra of McDonald data for BD+56◦ 0501.
267 Figure A-7: Stacked spectra of McDonald data for BD+56◦ 0563.
268 Figure A-8: Stacked spectra of McDonald data toward HD 14443.
269 Figure A-9: Stacked spectra of McDonald data toward HD 14476.
270 Figure A-10: Stacked spectra of McDonald data toward BD+56◦ 0578.
271 Figure A-11: Stacked spectra of McDonald data toward HD 14947.
272 Figure A-12: Stacked spectra of McDonald data toward HD 15629.
273 Figure A-13: Stacked spectra of McDonald data toward HD 16691.
274 Figure A-14: Stacked spectra of McDonald data toward HD 17505.
275 Figure A-15: Stacked spectra of McDonald data toward HD 17520.
276 Figure A-16: Stacked spectra of McDonald data for BD+60◦ 0586.
277 Figure A-17: Stacked spectra of McDonald data toward HD 19243.
278 Figure A-18: Stacked spectra of McDonald data toward HD 19820.
279 Figure A-19: Stacked spectra of McDonald data toward HD 25443.
280 Figure A-20: Stacked spectra of McDonald data toward HD 25638.
281 Figure A-21: Stacked spectra of McDonald data toward HD 45314.
282 Figure A-22: Stacked spectra of McDonald data toward HD 46223.
283 Figure A-23: Stacked spectra of McDonald data toward HD 192303.
284 Figure A-24: Stacked spectra of McDonald data toward HD 192641.
285 Figure A-25: Stacked spectra of McDonald data toward HD 193793.
286 Figure A-26: Stacked spectra of McDonald data toward HD 214449.
287 Figure A-27: Stacked spectra of McDonald data toward HD 217086.
288 Appendix B
H2 Spectra from FUSE
The H2 spectra shown here come from FUSE observations and are described in Section 3.4.3. Figures 3-4, 3-5, and 3-6 provide additional details about the plots of the H2 spectra.
289 HD 5689
1.00 iueB1 H B-1: Figure
0.75 2 pcrmtwr D5689. HD toward spectrum 290 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 12882
1.00 iueB2 H B-2: Figure
0.75 2 pcrmtwr D12882. HD toward spectrum 291 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 13841
1.00 iueB3 H B-3: Figure
0.75 2 pcrmtwr D13841. HD toward spectrum 292 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 13969
1.00 iueB4 H B-4: Figure
0.75 2 pcrmtwr D13969. HD toward spectrum 293 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 BD+56 0501
iueB5 H B-5: Figure 1.00
0.75 2 pcrmtwr BD+56 toward spectrum 294 0.50
0.25 ◦ 0501.
0.00
1040 1050 1060 1070 1080 1090 1100 BD+56 0563
iueB6 H B-6: Figure 1.00
0.75 2 pcrmtwr BD+56 toward spectrum 295 0.50
0.25 ◦ 0563.
0.00
1040 1050 1060 1070 1080 1090 1100 HD 14053
1.00 iueB7 H B-7: Figure
0.75 2 pcrmtwr D14053. HD toward spectrum 296 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 14443
1.00 iueB8 H B-8: Figure
0.75 2 pcrmtwr D14443. HD toward spectrum 297 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 14476
1.00 iueB9 H B-9: Figure
0.75 2 pcrmtwr D14476. HD toward spectrum 298 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 BD+56 0578
iueB1:H B-10: Figure 1.00
0.75 2 pcrmtwr BD+56 toward spectrum 299 0.50
0.25 ◦ 0578.
0.00
1040 1050 1060 1070 1080 1090 1100 HD 14947
1.00 iueB1:H B-11: Figure
0.75 2 pcrmtwr D14947. HD toward spectrum 300 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 iueB1:H B-12: Figure
HD 15629
1.00 hw eas fisrmn rbesi h bevto of observation the in bands. two problems other instrument the of because shown 2 pcrmtwr D169 nytobnso H of bands two Only 15629. HD toward spectrum
0.75 301 0.50
0.25
0.00
2 1040 1050 1060 1070 1080 1090 1100 are HD 16691
1.00 iueB1:H B-13: Figure
0.75 2 pcrmtwr D16691. HD toward spectrum 302 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 iueB1:H B-14: Figure
HD 17505
1.00 hw eas fisrmn rbesi h bevto of observation the in band. problems other instrument the of because shown 2 pcrmtwr D155 nytrebnso H of bands three Only 17505. HD toward spectrum
0.75 303 0.50
0.25
0.00
2 1040 1050 1060 1070 1080 1090 1100 are HD 17520
1.00 iueB1:H B-15: Figure
0.75 2 pcrmtwr D17520. HD toward spectrum 304 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 BD+60 0586
iueB1:H B-16: Figure 1.00
0.75 2 pcrmtwr BD+60 toward spectrum 305 0.50
0.25 ◦ 0586.
0.00
1040 1050 1060 1070 1080 1090 1100 HD 19243
1.00 iueB1:H B-17: Figure
0.75 2 pcrmtwr D19243. HD toward spectrum 306 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 19820
1.00 iueB1:H B-18: Figure
0.75 2 pcrmtwr D19820. HD toward spectrum 307 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 25443
1.00 iueB1:H B-19: Figure
0.75 2 pcrmtwr D25443. HD toward spectrum 308 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 25638
1.00 iueB2:H B-20: Figure
0.75 2 pcrmtwr D25638. HD toward spectrum 309 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 iueB2:H B-21: Figure
HD 45314
1.00 hw eas fisrmn rbesi h bevto of observation the in bands. two problems other instrument the of because shown 2 pcrmtwr D434 nytobnso H of bands two Only 45314. HD toward spectrum
0.75 310 0.50
0.25
0.00
2 1040 1050 1060 1070 1080 1090 1100 are HD 46223
1.00 iueB2:H B-22: Figure
0.75 2 pcrmtwr D46223. HD toward spectrum 311 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 192303
1.00 iueB2:H B-23: Figure
0.75 2 pcrmtwr D192303. HD toward spectrum 312 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 193793
1.00 iueB2:H B-24: Figure
0.75 2 pcrmtwr D193793. HD toward spectrum 313 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 iueB2:H B-25: Figure
HD 214419
h te ad h tfo h te ad soeli nthe on overlaid is bands of other observation the from band. the fit missing in The band. problems other instrument the of because shown 1.00 2 pcrmtwr D241.Ol he ad fH of bands three Only 214419. HD toward spectrum
0.75 314 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 2 are HD 217086
1.00 iueB2:H B-26: Figure
0.75 2 pcrmtwr D217086. HD toward spectrum 315 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 92964
1.00 iueB2:H B-27: Figure
0.75 2 pcrmtwr D92964. HD toward spectrum 316 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 97253
1.00 iueB2:H B-28: Figure
0.75 2 pcrmtwr D97253. HD toward spectrum 317 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 108927
1.00 iueB2:H B-29: Figure
0.75 2 pcrmtwr D108927. HD toward spectrum 318 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 149404
1.00 iueB3:H B-30: Figure
0.75 2 pcrmtwr D149404. HD toward spectrum 319 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 151932
1.00 iueB3:H B-31: Figure
0.75 2 pcrmtwr D151932. HD toward spectrum 320 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 HD 168076
1.00 iueB3:H B-32: Figure
0.75 2 pcrmtwr D168076. HD toward spectrum 321 0.50
0.25
0.00
1040 1050 1060 1070 1080 1090 1100 Appendix C
CO Spectra from FUSE
The CO spectra shown here come from FUSE observations. Section 3.4.3 describes the FUSE CO spectra. See Figure 3-9 for additional details about FUSE CO figures.
322 Figure C-1: CO spectra from FUSE toward HD 5689.
323 Figure C-2: CO spectra from FUSE toward HD 12882.
324 Figure C-3: CO spectra from FUSE toward HD 13841.
325 Figure C-4: CO spectra from FUSE toward HD 13969.
326 Figure C-5: CO spectra from FUSE toward BD+56◦ 0501.
327 Figure C-6: CO spectra from FUSE toward HD 14476.
328 Figure C-7: CO spectra from FUSE toward HD 15629.
329 Figure C-8: CO spectra from FUSE toward HD 16691.
330 Figure C-9: CO spectra from FUSE toward HD 17505.
331 Figure C-10: CO spectra from FUSE toward HD 17520.
332 Figure C-11: CO spectra from FUSE toward BD+60◦ 0586.
333 Figure C-12: CO spectra from FUSE toward HD 19243.
334 Figure C-13: CO spectra from FUSE toward HD 19820.
335 Figure C-14: CO spectra from FUSE toward HD 25638.
336 Figure C-15: CO spectra from FUSE toward HD 45314.
337 Figure C-16: CO spectra from FUSE toward HD 46223.
338 Figure C-17: CO spectra from FUSE toward HD 97253.
339 Figure C-18: CO spectra from FUSE toward HD 149404.
340 Figure C-19: CO spectra from FUSE toward HD 151932.
341 Figure C-20: CO spectra from FUSE toward HD 168076.
342 Figure C-21: CO spectra from FUSE toward HD 170740.
343 Figure C-22: CO spectra from FUSE toward HD 14053 and HD 14947. Only two bands of CO were measurable for each sight line.
344 Figure C-23: CO spectra from FUSE toward HD 193793 and HD 197770. Only two bands of CO were measurable for each sight line.
345 Figure C-24: CO spectra from FUSE toward HD 214419 and HD 217086. Only two bands of CO were measurable for each sight line.
346 Figure C-25: CO spectra from FUSE toward BD+56◦ 0563 and BD+56◦ 0578. Only two bands of CO were measurable for each sight line.
347 Figure C-26: CO spectra from FUSE toward HD 14443, HD 92964, HD 192303, and HD 192641. Only a single band of CO was mea- surable for each sight line.
348 Appendix D
CO Spectra from HST
The CO spectra shown here come from HST observations. Section 3.4.2 describes
the CO spectra and results from HST observations. Figures 3-7 and 3-8 provide addi-
tional details about high-resolution and medium-resolution HST spectra. Additional details for the figures with 13CO spectra are in the caption of Figure 3-11.
349 Figure D-1: CO spectra from HST toward HD 13841.
350 Figure D-2: CO spectra from HST toward HD 25443.
351 Figure D-3: 13CO spectra from HST toward HD 108927. See Figure 3-11 for additional details.
352 Figure D-4: CO spectra from HST toward HD 46223 (medium resolution).
353 Figure D-5: Spectrum of the d-X (5-0) CO band from HST toward HD 46223 (medium resolution). See Figure 3-12 for details.
354 Figure D-6: 13CO spectra from HST toward HD 46223 (medium resolution). See Figure 3-13 for details.
355 Figure D-7: CO spectra from HST toward HD 108927.
356 Figure D-8: CO spectra from HST toward HD 170740. See Figure 3-7 for additional details.
357 Appendix E
Spectra for GOT C+ Sight Lines
The spectra shown here come from the GOT C+ project and are described in
Section 3.5. See Figure 3-17 for additional details about the GOT C+ figures.
358 Figure E-1: (Top) Absorption spectra in Vlsr toward HD 168607. (Bottom) Emission spectra in Vlsr from the closest pointing, G014.8+1.0. The red ticks indicate the locations of fit components in all fig- ures in this Appendix.
359 Figure E-2: (Top) Absorption spectra in Vlsr toward HD 165918. (Bottom) Emission spectra in Vlsr from the closest pointing, G010.4+0.0.
360 Figure E-3: (Top) Absorption spectra in Vlsr toward HD 167498. (Bottom) Emission spectra in Vlsr from the closest pointing, G015.7+1.0.
361 Figure E-4: (Top) Absorption spectra in Vlsr toward HD 167812. (Bottom) Emission spectra in Vlsr from the closest pointing, G015.7+1.0.
362 Figure E-5: (Top) Absorption spectra in Vlsr toward HD 169754. (Bottom) Emission spectra in Vlsr from the closest pointing, G020.0+0.0.
363 Figure E-6: (Top) Absorption spectra in Vlsr toward HD 174509. (Bottom) Emission spectra in Vlsr from the closest pointing, G032.6+0.0.
364 ◦ Figure E-7: (Top) The absorption spectra in Vlsr toward BD+ 49 3482. (Bottom) Emission spectra in Vlsr from the closest pointing, G091.7+1.0.
365 ◦ Figure E-8: (Top) The absorption spectra in Vlsr toward BD+49 3484. (Bottom) Emission spectra in Vlsr from the closest pointing, G091.7+1.0.
366 Figure E-9: (Top) Absorption spectra in Vlsr toward HD 240179. (Bottom) Emission spectra in Vlsr from the closest pointing, G109.8+0.0.
367 Figure E-10: (Top) Absorption spectra in Vlsr toward HD 240183. (Bottom) Emission spectra in Vlsr from the closest pointing, G109.8+0.0.
368 Figure E-11: (Top) Absorption spectra in Vlsr toward HD 47073. (Bottom) Emission spectra in Vlsr from the closest pointing, G207.2+1.0.
369 Figure E-12: (Top) Absorption spectra in Vlsr toward HD 260737. (Bottom) Emission spectra in Vlsr from the closest pointing, G207.2+1.0.
370 Figure E-13: (Top) Absorption spectra in Vlsr toward HD 55469. (Bottom) Emission spectra in Vlsr from the closest pointing, G225.3+0.0.
371 Figure E-14: (Top) Absorption spectra in Vlsr toward HD 55981. (Bottom) Emission spectra in Vlsr from the closest pointing, G225.3+0.0.
372 Appendix F
Density and Temperature from C I
These figures shows the density of collisional partners versus temperature derived from the excitation levels of C I. The first panel with the star name shows the results for the total combined column density. The remainder of the panels show the results for each component. Within each panel, the solid lines correspond to the 3 sigma maximum and minimum density derived for the J =1/J =0 ratio and the dotted lines correspond to the 3 sigma maximum and minimum density derived for the J =2/J =0 ratio.
373 Figure F-1: This figure shows C I density of collisional partners versus tem- perature toward HD 13841.
374 Figure F-2: This figure shows C I density of collisional partners versus tem- perature toward HD 25443.
375 Figure F-3: This figure shows C I density of collisional partners versus tem- perature toward HD 46223.
376 Figure F-4: This figure shows C I density of collisional partners versus tem- perature toward HD 108927. This sight line only has one mean- ingful C I component.
377 Figure F-5: This figure shows C I density of collisional partners versus tem- perature toward HD 170740.
378