Filamentary dark clouds:structure, kinematics, and star formation:a study in Lupus and Norma

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Universidade de Lisboa

Faculdade de Ciˆ

encias

Departamento de F´

ısica

FILAMENTARY DARK CLOUDS:

STRUCTURE, KINEMATICS, AND

STAR FORMATION

A STUDY IN LUPUS AND NORMA

Miguel Centeno Moreira

Doutoramento em Astronomia e Astrof´

ısica

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FILAMENTARY DARK CLOUDS:

STRUCTURE, KINEMATICS, AND

STAR FORMATION

A STUDY IN LUPUS AND NORMA

Miguel Centeno Moreira

Doutoramento em Astronomia e Astrof´

ısica

Tese orientada pelo Prof. Doutor Jo˜ao Lin yun

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Acknowledgements

Over the years I have received generous help and advice from many individuals, colleagues, friends and family. It is an honour and a pleasure to acknowledge them. The work developed in this thesis was funded by the Funda¸c˜ao para a Ciˆencia e Tecnologia, under program PRAXIS XXI, BD/3451/94.

First, I should mention that I am profoundly indebted to Prof. Jo˜ao Lin Yun, my advisor, for his ability and tireless effort to motivate me toward the completion of both my MSc thesis and the present work toward a PhD degree in Astronomy & Astrophysics. Since 1992, he has continously opened for me the window of Astrophysical Research and all its wonders. We have also discovered together the incredible chilean landscape and night sky, from countless hours of observing time with several telescopes at ESO La Silla, and CTIO.

I am also extremely grateful to Dr. Daniele Galli for his advice, careful reading of this manuscript, and last but not least, his unmatched hospitality during my stay at Osservatorio Astrofisico di Arcetri, in Firenze, Italy. Among the Arcetri star formation group, I have also greatly benefited from discussions with Drs. Malcolm Walmsley and Francesco Palla, whom I so admire.

I would also like to thank, in particular, the SEST staff for providing such a wonderful and easy operation of this beautiful radiotelescope, and assisting with the setup of the observing runs.

Finally, I would like to thank my family with everlasting love, to whom I dedicate this piece of research: my parents, Yvette and Binau, Rosine, Lu´ısa (a.k.a. Bu), my brothers Berna, Pedro, and Junho, and my four beautiful children, Luisinha, Vasco, Rodrigo, and Bernardo. In the words of Carl Sagan, it is a priviledge and an honour to share a planet and an epoch with them all.

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Resumo em L´ıngua Portuguesa

Compreender como as estrelas se formam a partir do meio interestelar difuso à escala global e local é um poblema fundamental que permanece em aberto na astrof´ısica contem-porânea. Das nuvens difusas às nuvens moleculares, e da´ı para a forma¸cão, vida e morte de estrelas e planetas, o plano galáctico é o local onde todas as fases do ciclo de vida da Galáxia podem ser estudadas em contexto. Mais de mil nuvens escuras são vis´ıveis num raio de várias centenas de parsec em redor do Sol, e uma fraçcão substancial delas contém jovens estrelas de pequena massa. A grande variedade de tamanhos, formas, massas, e com-plexidade destas nuvens e suas estrelas vis´ıveis associadas sugerem que as nuvens escuras observáveis abrangem uma vasta gama de estágios evolutivos, e que estudos comparativos de nuvens escuras e do seu conteúdo estelar podem revelar pistas importantes sobre a f´ısica das nuvens e dos processos que conduzem à forma¸cão de estrelas. Muitos dos complexos de nuvens que formam estrelas possuem estruturas alongadas com tamanhos da ordem do par-sec, como demonstram imagens obtidas no óptico, infravermelho, e submil´ımetro, bem como dados obtidos nas transi¸cões rotacionais da molécula de CO. Várias destas nuvens exibem uma estrutura tipo cabe¸ca-cauda na qual uma componente central densa aproximadamente esférica está associada a uma cauda alongada menos densa. Observa¸cões recentes revelaram uma profusão de filamentos em nuvens moleculares galácticas com uma extensão da ordem do parsec que sugerem uma liga¸cão ´ıntima entre a estrutura filamentar do meio interestelar e o processo de forma¸cão de núcleos densos das nuvens moleculares.

O objectivo desta tese é o estudo comparativo das condi¸cões f´ısicas e actividade de forma¸cão de estrelas das nuvens escuras filamentares GF 17, GF 20, e Sa 187/188. Procu-ramos aprender sobre a forma¸cão, evolu¸cão e destrui¸cão das nuvens escuras filamentares, através do estudo das suas semelhanas e diferen¸cas, quer estruturais quer cinemáticas. Um tal estudo deverá fornecer pistas sobre a origem das nuvens filamentares, mas também permitir impor restri¸cões significativas nos modelos teóricos actualmente existentes sobre estabilidade e colapso de estruturas filamentares. Para este efeito, foram obtidos e analisados mapas de riscas moleculares em transi¸cões rotacionais de CO, 13_{CO, C}18_{O, e CS no mil´ımetro, e ainda}

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iii

Atlas, e imagens no infravermelho próximo obtidas com o telescó ESO de 2.2 m de diâmetro, situado em La Silla, no Chile. Com a estrutura e cinemática das nuvens filamentares como motiva¸cão cient´ıfica, podemos resumir o plano desta disserta¸cão da seguinte forma:

1) a partir de um modelo simples de excita¸cão da poeira, obter mapas da temperatura de cor, profundidade óptica, e extin¸cão visual da poeira interestelar;

2) a partir dos dados das riscas moleculares observadas, construir mapas detalhados da distribui¸cão do gás nas regiões centrais e filamentares de cada nuvem, com uma resolu¸cão es-pacial sem precedentes nestas frequências. Em particular, usar mapas de densidade de coluna do gás para os relacionar com os mapas da poeira, e por outro lado, para construir perfis radiais para comparar¸cão com modelos teóricos de distribui¸cão de densidade em núcleos densos esféricos, e em estruturas filamentares com simetria cil´ındrica;

3) usar a informa¸cão cinemática do gás para inferir a dinâmica das nuvens, e investigar a presen¸ca de gradientes de velocidade ao longo da sua estrutura;

4) usar os dados de rádio e infravermelho próximo para identificar eventuais novos sinais de actividade de forma¸cão de estrelas nestas nuvens.

Come¸camos esta investiga¸cão no Cap´ıtulo 3 onde descrevemos e analisamos as pro-priedades globais da poeira nas nuvens GF 17, GF 20 e Sa 187/188, nos comprimentos de onda de 60 e 100 µm. Exploramos uma compara¸cão entre a emissão térmica da poeira e a emissão do gás molecular detectado. Para esta compara¸cão foi necessário degradar a reso-lu¸cão angular dos mapas da emissão do gás molecular para a resoreso-lu¸cão angular das imagens IRAS, que é pior. Conclui-se que existe uma boa correla¸cão entre a profundidade óptica da poeira a 100µm e a intensidade integrada de 13_{CO em GF 17 e GF 20, o que significa que a}

emissão da poeira no infravermelho long´ınquo pode ser usada como indicador da densidade de coluna do gás molecular nestas nuvens. Contudo, tal não acontece no caso de Sa 187/188, devido à presen¸ca de vários objectos estelares jovens. Verificamos que a popula¸cão de objec-tos estelares jovens de GF 17 e GF 20 não consegue explicar a luminosidade destas nuvens entre 60 e 100µm, uma vez que a sua contribui¸cão é de apenas 20%. Conclu´ımos que a fonte dominante de aquecimento da poeira nestas nuvens tem que ser externa, sendo provável que

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seja a radia¸c˜ao proveniente do subgrupo Upper-Scorpius da associa¸c˜ao de estrelas Scorpius OB2.

Prosseguimos a nossa investiga¸cão no Cap´ıtulo 4, no qual fazemos uma descri¸cão dos modelos de esferas e cilindros gasosos com equa¸cão de estado politrópica, com os quais modelamos respectivamente os perfis radiais de densidade de coluna dos núcleos densos e das regiões filamentares, obtidos a partir dos mapas das riscas moleculares.

No Cap´ıtulo 5 caracterizamos as propriedades f´ısicas dos núcleos densos, com as quais aferimos o seu estado dinâmico através de uma análise do balan¸co energético. Identificamos 9 núcleos densos em GF 17, 8 núcleos em GF 20 e 6 núcleos em Sa 187/188. Todos os núcleos (com excep¸cão de GF17-C6 e GF17-C7) foram mapeados em mais do que uma transi¸cão molecular. Os núcleos densos apresentam semelhan¸ca significativa de linha para linha no seu alongamento relativo, tamanho e orienta¸cão, suportando a ideia de que os ma-pas de 13_{CO, C}18_{O, e CS correspondem essencialmente ao mesmo volume de gás. Obtemos}

valores de pico de densidade de coluna de H2 no intervalo 4.1 − 10.9 × 1021cm−2 em GF 17 e

GF 20, e valores mais elevados em Sa 187/188, no intervalo 8.5 − 13.5 × 1021_cm−2_{. O n´}_ucleo

denso t´ıpico desta amostra possui uma distribui¸cão alongada da sua densidade de coluna, com um rácio axial projectado no intervalo 0.50 − 0.60, em escalas de tamanho de 0.07 pc para os núcleos de GF 17 e GF 20, e 0.15 pc para os núcleos de Sa 187/188. Estes núcleos são carac-terizados por larguras de risca moderadamente supersónicas, implicando que são suportados por mecanismos não térmicos. Neste contexto, a natureza não térmica das rela¸cões largura de risca–tamanho investigadas favorecem a turbulência (e não a gravidade) como mecanismo responsável. Contudo, estas rela¸cões apresentam uma diferen¸ca significativa com o resultado clássico, no sentido em que os núcleos densos do presente estudo exibem larguras de risca sistematicamente maiores. Esta diferen¸ca, igualmente verificada em glóbulos de Bok, suge-re uma semelhan¸ca cinemática que pode significar uma liga¸cão evolutiva entsuge-re filamentos globulares e glóbulos de Bok. Este aspecto merece uma investiga¸cão cuidada no futuro.

Os gradientes de velocidade encontrados (tipicamente 0.2 − 1.6 km s−1 _pc−1_{) mostram}

que a rota¸cão dos núcleos, a existir de todo, não é relevante em termos energéticos. Através de uma análise virial dos núcleos, verificamos que a maioria dos núcleos são provavelmente

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v

estruturas estabilizadas por uma pressão externa. Esta conclusão é totalmente compat´ıvel com a modela¸cão dos perfis de densidade, que indica que os modelos que melhor se ajustam correspondem a estruturas estabilizadas e truncadas por pressão externa, caracterizadas por ´ındices politrópicos baixos (N ≤ 2), e portanto não isotérmicos.

No Cap´ıtulo 6 exploramos as propriedades f´ısicas e dinâmicas das nuvens filamentares como um todo, fazendo uma caracteriza¸cão detalhada da sua estrutura de densidade e ainda da distribui¸cão em larga escala da velocidade do gás. Esta última revela a existência de movi-mentos ao longo de toda a estrutura das nuvens com uma magnitude de ∼ 0.5 km s−1 _pc−1_.

Estes gradientes de velocidade não podem ser interpretados como simples rota¸cão ou efeitos de maré galáctica. Em particular, argumentamos que no caso de GF 17 e GF 20 estes gradientes são uma indica¸cão muito forte de que as nuvens foram atravessadas por uma frente de choque em propaga¸cão, provocada pelas camadas em expansão Upper-Scorpius e Upper-Centaurus-Lupus. No caso de Sa 187/188, a origem não é clara, embora o gradiente de velocidade encontrado seja compat´ıvel com o cenário no qual movimentos de larga escala do meio interestelar (provocados por uma explosão de supernova) conduzem à forma¸cão de uma estrutura el´ıptica achatada (tipo folha) que colapsou em cerca de 2 × 106 _anos.

A análise virial das nuvens, e em particular das suas regiões filamentares, revela mais uma vez que a rota¸cão não é relevante em termos energéticos. Revela ainda que as regiões centrais de GF 17 e GF 20 possuem energias cinética e gravitacional comparáveis, enquanto em Sa 187/188 a gravidade predomina por um factor de 20%. Contudo, as regiões fila-mentares das três nuvens revelam uma semelhan¸ca surpreendente, que consiste no facto da energia cinética ser 3 vezes superior à energia gravitacional, o que sugere que estas regiões se encontram em estados dinâmicos muito semelhantes. As velocidades do gás são supersónicas nas três nuvens, e portanto o mecanismo de suporte é essencialmente não térmico. Contudo, as distribui¸cões de velocidade obtidas para 13_{CO, C}18_{O, e CS parecem ser intrinsecamente}

diferentes, o que pode significar que o gás mais denso possui uma dispersão de velocidades menor do que as regiões entre núcleos densos. Isto sugere que o mecanismo de forma¸cão dos núcleos densos foi acompanhada de uma certa dissipa¸cão dos movimentos não térmicos, mas a prevalência destes movimentos sugere a necessidade de uma pressão externa para

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estabi-lizar as nuvens, impedindo que estas se dispersem no meio interestelar. A modela¸cão das estruturas filamentares em GF 17 e GF 20 revela que estas são bem descritas por perfis de densidade com ´ındice politrópico baixo (respectivamente, N = +1 e N = +0.5) truncados por pressão externa, cuja densidade central é inferior a 4 × 104 _cm−3_{. O n´}_{umero de n´}_ucleos

densos observados em cada região filamentar, bem como a sua separa¸cão espacial, parecem ser melhor explicados pela fragmenta¸cão de estruturas cil´ındicas com ´ındices politrópicos baixos do que de estruturas isotérmicas (N=∞). A continuidade observada no padrão de velocidades ao longo de toda a estrutura das nuvens GF 17 e GF 20 implica que o processo de forma¸cão de núcleos densos nestas nuvens não desacoplou o gás denso do gás menos denso na sua vizinhan¸ca.

Finalmente, no Cap´ıtulo 7 mostramos que é muito provável que a estrela T Tauri RU Lupi seja responsável por uma ejeçcão de material do tipo “outflow” bipolar, bem colimado e praticamente alinhado com a linha de visão para a nuvem GF 20. Estimamos que a massa deste “outflow” seja de ∼ 3.3 × 10−3 _M

, e que a sua luminosidade mecˆanica seja de

∼ 5.2 × 10−5 _L

. A idade dinˆamica do “outflow” ´e estimada em ∼ 2 × 104 anos, e a taxa

de perda de massa ´e avaliada em 1.7 × 10−7_M

ano−1. Por outro lado, mostramos que a

fonte IRAS 16295−4452 associda ao n´ucleo denso Sa187-C3 ´e uma estrela jovem de Classe I, com uma luminosidade total de 35 L. O perfil das riscas de CO em torno desta fonte

mostra a presen¸ca de um “outflow” que pode ser originado pela fonte IRAS 16295−4452. Os nossos resultados apontam igualmente para a presen¸ca de um “outflow” bipolar no núcleo denso Sa187-C2, originado por uma fonte embebida não detectada. A massa contida nos dois lóbulos é ≥ 0.011 M, e o momento linear e a energia cinética do “outflow” são estimados

em 0.07 M kms−1 e 0.2 M km2s−2, respectivamente. A luminosidade mecˆanica ´e de

3 × 10−4 _L

, enquanto idade dinâmica é avaliada em ∼ 105 anos. Estes parâmetros estão

de acordo com os dados dinˆamicos do “outflow” molecular HH 56/57 localizado na mesma regi˜ao central de Sa 187/188.

PALAVRAS-CHAVE: Meio Interestelar: Nuvens – Meio Interestelar: Evolu¸c˜ao – Nuvens: Individual (Lupus 2 , Lupus 4, Sa 187/188) – Nuvens: Estrutura – Estrelas: Forma¸c˜ao

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vii

Abstract

Recent observations have revealed a profusion of parsec-scale filaments in galactic mole-cular clouds and suggest an intimate connection between the filamentary structure of the ISM and the process of formation of dense cloud cores. The goal of this thesis is the comparative study of the physical conditions and star formation activity within the filamentary dark clouds GF 17, GF 20, and Sa 187/188. We seek to learn about the formation, evolution, and destruction of filamentary dark clouds by studying the structural and kinematical similarities and differences between them. Millimeter molecular line maps in rotational transitions of CO, 13_{CO, C}18_{O, and CS, and IRAS Sky Survey Atlas images at 12, 25, 60, and 100 µm}

were analysed.

We identify 9 cores in GF 17, 8 cores in GF 20 and 6 cores in Sa 187/188. The dense cores exhibit significant similarity from line to line in their relative elongation, size, and orientation, supporting the idea that the 13_{CO, C}18_{O, and CS line maps are largely sampling the same}

volume of gas. Filament and dense core kinematics are both dominated by nonthermal supersonic internal motions, suggesting that core formation has not resulted in kinematical decoupling of the dense gas from its parental cloud. The filamentary regions have radial column density profiles with a central flattened region and a tail well fitted by pressure truncated polytropic cylinders with small polytropic indices. The picture that emerges is that some form of cylindrical supersonic fragmentation has played a role in the formation of the dense cores contained in these clouds, as the filamentary gas prior to fragmentation has still a supersonic level of turbulence which was not dissipated when the filament condensed out of the ISM. The age spread of the YSOs in these clouds and the large-scale velocity gradients found are consistent with the star formation activity having been triggered by an external event that shaped these filamentary clouds.

KEYWORDS: ISM: clouds – ISM: evolution – Clouds: individual (Lupus 2 , Lupus 4, Sa 187/188) – Clouds: structure – Stars: formation

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1 Introduction 1

1.1 The Filamentary Nature of the Interstellar Medium . . . 1

1.2 The Origin of Filamentary Cloud Structures . . . 2

1.3 Filamentary Dark Clouds and Star Formation . . . 3

1.4 Aim of the Thesis . . . 7

1.5 Overview of the Subject Clouds . . . 10

1.5.1 Selection Criteria . . . 10

1.5.2 GF 17 and GF 20 . . . 11

1.5.3 Sa 187/188 . . . 14

1.6 Thesis Plan . . . 16

2 Millimeter Molecular Line Mapping 17 2.1 Introduction . . . 17

2.2 Observations and Data Reduction . . . 19

2.2.1 Observations . . . 19

2.2.2 Data Reduction . . . 20

2.3 Maps . . . 24

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CONTENTS ix 2.4 LTE Analysis . . . 58 2.4.1 Excitation Temperature . . . 58 2.4.2 Optical Depth . . . 59 2.4.3 Column Densities . . . 61 2.4.4 Mass . . . 63

2.5 Molecular Gas Structure in GF 17, GF 20, and Sa 187/188 . . . 64

2.6 Identification of Dense Cores . . . 65

2.7 Spectral Line Profiles . . . 67

2.7.1 GF 17 . . . 67 2.7.2 GF 20 . . . 72 2.7.3 Sa 187/188 . . . 74 2.7.4 BHR 86 . . . 76 2.7.5 CB 63 . . . 77 2.8 Summary of Chapter 2 . . . 79

3 Dust Emission in GF 17, GF 20, and Sa 187/188 80 3.1 Introduction . . . 80

3.2 Data Acquisition and Image Processing . . . 82

3.3 Obtaining Images of Dust Temperature and Optical Depth . . . 83

3.3.1 Dust Temperature . . . 83

3.3.2 Dust Optical Depth and Visual Extinction . . . 84

3.3.3 Errors and Assumptions . . . 85

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3.4.1 IRAS Cloud Morphology . . . 92

3.4.2 Dust Color Temperature and Optical Depth . . . 93

3.4.3 Dust Mass . . . 97

3.4.4 Correlations Between Gas and Dust . . . 99

3.4.5 Visual Extinction . . . 105

3.5 Evidence for Smooth Cloud Edges in GF 17 and GF 20 . . . 108

3.6 Comparison with Infrared Cirrus Clouds and Molecular Dark Clouds . . . . 110

3.7 Infrared Luminosity-to-Mass Ratio . . . 113

3.8 The Source of Dust Heating in GF 17 and GF 20 . . . 114

3.9 Summary of Chapter 3 . . . 116

4 Density Structure of Self-gravitating Polytropes 118 4.1 Introduction . . . 118

4.2 Equilibrium of Self-gravitating Polytropes . . . 122

4.3 Physical Characteristics . . . 126

4.4 _{The Isothermal Case (N = ±∞) . . . 127}

4.5 Radial Profiles . . . 128

4.6 Obtaining Column Density Profiles . . . 129

4.6.1 Inclination to the Line of Sight . . . 132

4.7 Obtaining Best-fit Parameters . . . 134

5 The Properties of the Dense Cores 135 5.1 Introduction . . . 135

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CONTENTS xi

5.2 Results and Analysis . . . 137

5.2.1 Emission Strength . . . 137

5.2.2 Comparison With Surveys of Dense Cores in Dark Cloud Complexes . 140 5.2.3 Comparison With Surveys of Isolated Bok Globules . . . 145

5.2.4 Column Densities Derived from 13_{CO and C}18_{O Observations. . . 150}

5.2.5 Core Shapes and Sizes . . . 152

5.2.6 Velocity Gradients . . . 160

5.2.7 Line Width – Size Relations . . . 168

5.3 The Nonthermal Nature of the Line Width – Size Relations . . . 175

5.4 Transition to Coherence? . . . 178

5.5 Comparison with Bok Globules . . . 180

5.6 Energy Balance and Stability of the Cores . . . 182

5.7 Radial Column Density Profiles . . . 191

6 Kinematics and Structure of GF 17, GF 20, and Sa 187/188 204 6.1 Introduction . . . 204

6.2 Results . . . 205

6.2.1 Velocity Structure of GF 17 and GF 20 . . . 205

6.2.2 Velocity Structure of Sa 187/188 . . . 210

6.2.3 Cloud Energy Balance and Stability . . . 213

6.3 Density Structure of the Filamentary Regions . . . 217

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6.3.2 Comparison with Models: the Singular Logatrope . . . 222

6.3.3 Comparison with Models: the Polytropic Cylinder with Finite Index N 223 6.4 Spatial Distribution of nonthermal Motions and Consequences for Turbulent Models of Core Formation . . . 227

6.5 Continuity of the Velocity Field . . . 232

6.6 Cloud Fragmentation . . . 239

6.7 The Origin of the Cloud Structures in GF 17 and GF 20 . . . 244

6.8 The Origin of the Cloud Structure in Sa 187/Sa 188 . . . 248

7 Signs of Star Formation Activity in GF 20 and Sa 187/188 254 7.1 Introduction . . . 254

7.2 Is RU Lupi Driving a Molecular Outflow? . . . 254

7.3 _{IRAS 16295−4452: a Class I source in Sa187-C3 . . . 259}

7.3.1 Near-infrared Observations and Data Reduction . . . 261

7.3.2 _{The Molecular Environment of IRAS 16295−4452 . . . 263}

7.3.3 Spectral Line Profiles and Gas Kinematics . . . 266

7.3.4 _{IRAS 16295−4452: A Class I Protostar . . . 267}

7.3.5 Additional Near-infrared Sources Near the Dense Core . . . 271

7.4 Sa187-C2: A Star Forming Core? . . . 273

7.5 Sa187-C6: A Contracting or Pulsating Turbulent Core? . . . 275

7.6 Star Formation in Sa 187/188 . . . 278

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CONTENTS xiii

8 Conclusions and Future Prospects 281

8.1 The Structure of Globular Filaments . . . 281

8.2 Globular Filament Cores and Bok Globules . . . 283

8.3 Future Prospects . . . 284 8.3.1 Work in Progress . . . 285 Bibliography . . . 287 A Map Positions 300 A.1 GF 17 . . . 300 A.2 GF 20 . . . 302 A.3 Sa 187/188 . . . 305

A.4 Bok globules BHR 86 and CB 63 . . . 308

A.4.1 BHR 86 . . . 308

A.4.2 CB 63 . . . 309

B Spectral Line Maps 312 B.1 GF 17 . . . 312

B.2 GF 20 . . . 322

B.3 Sa 187/188 . . . 332

B.4 Bok globules BHR 86 and CB 63 . . . 338

B.4.1 BHR 86 . . . 338

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1.1 Digital Sky Survey images of the globular filaments GF 17 and GF 20 . . . . 13 1.2 Digital Sky Survey image of the filamentary dark cloud Sa 187/188 . . . 15

2.1 Spectral line profiles. . . 21 2.2 Map of CO excitation temperature toward the filamentary region of GF 17. . 26 2.3 Map of CO excitation temperature toward the main core region of GF 17. . . 26 2.4 Map of13_{CO (1−0) integrated emission toward the filamentary region of GF 17. 27}

2.5 Map of13_{CO (1 − 0) integrated emission toward the main core region of GF 17. 27}

2.6 Map of 13_{CO (1 − 0) antenna temperature toward the filamentary region of}

GF 17. . . 28 2.7 Map of13_{CO (1−0) antenna temperature toward the main core region of GF 17. 28}

2.8 Map of13_{CO (1 − 0) line velocity toward the filamentary region of GF 17. .} ₂₉

2.9 Map of13_{CO (1 − 0) line velocity toward the main core region of GF 17. . .} ₂₉

2.10 Maps of C18_{O (1 − 0) integrated emission and antenna temperature toward}

the filamentary region of GF 17. . . 30 2.11 Map of C18_{O (1 − 0) integrated emission toward the main core region of GF 17. 30}

2.12 Map of C18_{O (1−0) antenna temperature toward the main core region of GF 17. 31}

2.13 Map of C18_{O (2 − 1) integrated emission toward the main core region of GF 17. 31}

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LIST OF FIGURES xv

2.14 Map of C18_{O (2−1) antenna temperature toward the main core region of GF 17. 32}

2.15 Map of CS (2 − 1) velocity integrated emission toward the main core region

of GF 17. . . 33

2.16 Map of CS (2 − 1) antenna temperature toward the main core region of GF 17. 33 2.17 Map of CS (3 − 2) velocity integrated emission toward the main core region of GF 17. . . 34

2.18 Map of CS (3 − 2) antenna temperature toward the main core region of GF 17. 34 2.19 Map of CO excitation temperature toward GF 20. . . 36

2.20 Map of 13_{CO (1 − 0) integrated emission toward GF 20. . . .} ₃₆

2.21 Map of 13_{CO (1 − 0) antenna temperature toward GF 20. . . .} ₃₇

2.22 Map of 13_{CO (1 − 0) line velocity toward GF 20. . . .} ₃₇

2.23 Map of C18_{O (1−0) integrated emission toward the filamentary region of GF 20. 38} 2.24 Map of C18_{O (1 − 0) antenna temperature toward the filamentary region of} GF 20. . . 38

2.25 Map of CS (2 − 1) integrated emission toward the filamentary region of GF 20. 39 2.26 Map of CS (2 − 1) integrated emission toward the main core region of GF 20. 39 2.27 Map of CS (2 − 1) antenna temperature toward the filamentary region of GF 20. 40 2.28 Map of CS (2 − 1) antenna temperature toward the main core region of GF 20. 40 2.29 Map of CO excitation temperature toward the filamentary region (Sa 188). . 42

2.30 Map of CO excitation temperature toward the main core region (Sa 187). . . 42 2.31 Map of13_{CO (1 −0) integrated emission toward the filamentary region (Sa 188). 43}

2.32 Map of 13_{CO (1 − 0) integrated emission toward the main core region (Sa 187). 43}

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2.34 Map of13_{CO (1 −0) antenna temperature toward the main core region (Sa 187). 44}

2.35 Map of 13_{CO (1 − 0) line velocity toward the filamentary region (Sa 188). . .} ₄₅

2.36 Map of 13_{CO (1 − 0) line velocity toward the main core region (Sa 187). . . .} ₄₅

2.37 Map of CS (3 − 2) integrated emission toward the filamentary region (Sa 188). 46 2.38 Map of CS (3 − 2) integrated emission toward the main core region (Sa 187). 46 2.39 Map of CS (3 − 2) antenna temperature toward the filamentary region (Sa 188). 47 2.40 Map of CS (3 − 2) antenna temperature toward the main core region (Sa 187). 47 2.41 Map of CO excitation temperature toward BHR 86. . . 49

2.42 Map of 13_{CO (1 − 0) integrated emission toward BHR 86. . . .} ₄₉

2.43 Map of 13_{CO (1 − 0) antenna temperature toward BHR 86. . . .} ₅₀

2.44 Map of C18_{O (2 − 1) integrated emission toward BHR 86. . . .} ₅₀

2.45 Map of C18_{O (2 − 1) antenna temperature toward BHR 86. . . .} ₅₁

2.46 Map of CS (2 − 1) integrated emission toward BHR 86. . . 51

2.47 Map of CS (2 − 1) antenna temperature toward BHR 86. . . 52

2.48 Map of CS (3 − 2) integrated emission toward BHR 86. . . 52

2.49 Map of CS (3 − 2) antenna temperature toward BHR 86. . . 53

2.50 Map of CO excitation temperature toward CB 63. . . 55

2.51 Map of 13_{CO (1 − 0) integrated emission toward CB 63. . . .} ₅₅

2.52 Map of 13_{CO (1 − 0) antenna temperature toward CB 63. . . .} ₅₆

2.53 Map of 13_{CO (1 − 0) line velocity toward CB 63. . . .} ₅₆

2.54 Map of CS (2 − 1) integrated emission toward CB 63. . . 57

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LIST OF FIGURES xvii

2.56 Distribution of CO excitation temperature and 13_{CO (1 − 0) optical depth. .} ₆₀

2.57 Location of the dense cores in GF 17 . . . 68

2.58 Location of the dense cores in GF 20 . . . 69

2.59 Location of the dense cores in Sa 187/188 . . . 70

2.60 Line profiles in GF 17 . . . 71

2.61 Line profiles in the main core region of GF 20 . . . 73

2.62 Line profiles toward Sa 187 . . . 74

2.63 Line profiles toward Sa 187-C2 and Sa187-C6 . . . 75

2.64 Line profiles toward BHR 86 . . . 77

2.65 Line profiles toward CB 63 . . . 78

3.1 GF 17: IRAS co-added images . . . 89

3.2 GF 20: IRAS co-added images . . . 90

3.3 Sa 187/188: IRAS co-added images . . . 91

3.4 GF 17: temperature and opacity images . . . 95

3.5 GF 20: temperature and opacity images . . . 96

3.6 Sa 187/188: temperature and opacity images . . . 98

3.7 100 µm optical depth versus 13_{CO integrated emission for GF 17 and GF 20. 100} 3.8 100 µm optical depth versus C18_{O integrated emission for GF 17 and GF 20. 101} 3.9 100 µm optical depth versus 13_{CO integrated emission for Sa 187/188. . . 103}

3.10 Dust temperature versus gas column density for GF 17 and GF 20. . . 104

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3.12 The σ − AV relations with spatial filters of differing angular resolution for

GF 17 and GF 20 . . . 109

3.13 100 µm intensity versus molecular hidrogen column density. . . 112

4.1 Solutions of the Lane-Emden equations. . . 123

4.2 Radial profiles of the gas temperature, pressure, volume density, and mass for cylindrical polytropes . . . 130

4.3 Radial profiles of the gas temperature, pressure, volume density, and mass for spherical polytropes . . . 130

4.4 Line-of-sight integration of the density profiles . . . 131

4.5 Beam-convolved H2 column density profile of a self-gravitating cylinder of polytropic index N = 5. . . 132

5.1 Comparison with the Vilas-Boas, Myers, & Fuller (1994) and Vilas-Boas, My-ers, & Fuller (2000) surveys of dense cores. . . 143

5.2 Survey results for starless Bok globules. . . 147

5.3 _{Comparison with Bok globules: CS (2 − 1) data . . . 148}

5.4 _{Comparison with Bok globules: CS (3 − 2) data . . . 149}

5.5 _{Maps of CS (2 − 1) integrated emission towards Sa187-C1, Sa187-C2, and} Sa187-C3. . . 153

5.6 Maps of C18_{O (2 − 1) and CS (2 − 1) integrated emission toward Sa187-C4. . 154}

5.7 Half-maximum contour maps. . . 157

5.8 Number distribution of the difference in long axis position angle. . . 159

5.9 _{Comparison between CS (2 − 1) and C}18_{O emission . . . 160}

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LIST OF FIGURES xix

5.11 Comparisons of gradient magnitude and direction among different tracers . . 167

5.12 line width – size relations from CS data. . . 170

5.13 line width – size relations in GF 17 and GF 20. . . 173

5.14 line width – size relations for all tracers. . . 174

5.15 Thermal and nonthermal components of the line width – size relation. . . 177

5.16 Line width – size relations for a sample of Bok globules . . . 181

5.17 Dense core virial equilibrium analysis. . . 185

5.18 Critical external pressure and mean density. . . 190

5.19 Column density profiles of the dense cores in GF 17 . . . 192

5.20 Column density profiles of the dense cores in GF 20. . . 193

5.21 Column density profiles of the dense cores in Sa 187. . . 194

5.22 Column density profiles of the dense cores Sa187-C5 and Sa187C6. . . 195

5.23 Comparison with polytropic spheres in GF 17 . . . 199

5.24 Comparison with polytropic spheres in GF 20 . . . 200

6.1 Variation of the LSR velocity across GF 17 and GF 20. . . 206

6.2 Variation of the LSR velocity across Sa 187/188. . . 211

6.3 Distribution of the nonthermal velocity dispersion. . . 216

6.4 Column density profile of the filaments . . . 219

6.5 Comparison with polytropic filaments . . . 225

6.6 Distribution of the nonthermal velocity dispersion as a function of position. . 229

6.7 nonthermal velocity dispersion as a function of radius. . . 231

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6.9 Velocity oscillations in the GF 17 filament. . . 238 6.10 Global distribution of the H I column density around Ophiuchus and Lupus . 245 6.11 Distribution of observed radial velocities in GF 17 and GF 20 versus galactic

latitude. . . 247

7.1 RU Lupi and the main core region of GF 20. . . 255 7.2 _{Map of CO (1 − 0) wing integrated emission toward the main core of GF 20} 257 7.3 Young Stellar Objects known in Sa 187/188 . . . 260 7.4 Uncertainties and histograms of the J, H, and K magnitudes . . . 262 7.5 _{Maps of CO (1 − 0),}13_{CO (1 − 0) and CS (3 − 2) integrated emission toward}

IRAS 16295−4452. . . 263 7.6 _{HCN (1 − 0) spectra taken toward IRAS 16295−4452. . . 265} 7.7 _{The outflow associated with IRAS 16295−4452. . . 266} 7.8 _{Identification of the near-infrared counterpart of IRAS 16295−4452, Sa 187 IRS.268} 7.9 _{Spectral Energy Distribution of IRAS 16295−4452. . . 269} 7.10 Histograms of the H − K colors of all the stars detected at both H and K

near IRAS 16295−4452 . . . 272 7.11 CO and 13_{CO line profiles toward Sa187-C2 . . . 274}

7.12 2MASS Ks image of the region near Sa187-C2 . . . 275

7.13 Spectral energy distribution of G338.5374+02.0066 . . . 276 7.14 CS (2 − 1) line profiles toward Sa187-C6 . . . 277 8.1 Near-infrared mosaic of the globular filament GF 20 . . . 286

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LIST OF FIGURES xxi

A.2 GF 17: C18_{O (2 − 1) and CS (2 − 1) map positions . . . 302}

A.3 GF 20: CO (1 − 0) and 13_{CO (1 − 0) map positions . . . 303}

A.4 GF 20: C18_{O (1 − 0) and CS (2 − 1) map positions . . . 304}

A.5 Sa 187/188: CO (1 − 0) map positions . . . 305 A.6 Sa 187/188: 13_{CO (1 − 0) and C}18_{O (2 − 1) map positions . . . 306}

A.7 Sa 187/188: CS (2 − 1) and CS (3 − 2) map positions . . . 307 A.8 BHR 86: CO (1 − 0), 13_{CO (1 − 0) and C}18_{O map positions . . . 308}

A.9 BHR 86: CS (2 − 1), and CS (3 − 2) map positions . . . 309 A.10 CB 63: CO (1 − 0) map positions . . . 310 A.11 CB 63: CS (2 − 1) map positions . . . 311 B.1 GF 17: CO (1 − 0) spectra . . . 313 B.2 GF 17: 13_{CO (1 − 0) spectra . . . 314} B.3 GF 17: C18_{O (1 − 0) spectra . . . 315} B.4 GF 17: C18_{O (2 − 1) and CS (2 − 1) spectra . . . 316} B.5 GF 17: CS (3 − 2) spectra . . . 317 B.6 GF17-C1 and GF17-C2: CO (1 − 0) spectra . . . 318 B.7 GF17-C3 and GF17-C4: CO (1 − 0) spectra . . . 319 B.8 GF17-C1 and GF17-C2: 13_{CO (1 − 0) spectra . . . 320} B.9 GF17-C3 and GF17-C4: 13_{CO (1 − 0) spectra . . . 321}

B.10 GF 20: CO (1 − 0) and 13_{CO (1 − 0) spectra toward the main core region . . 323}

B.11 GF 20: CO (1 − 0) spectra toward the filamentary region . . . 324 B.12 GF 20: 13_{CO (1 − 0) spectra toward the filamentary region . . . 325}

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B.13 GF 20: C18_{O (1 − 0) spectra toward the filamentary region . . . 326} B.14 GF 20: CS (2 − 1) spectra . . . 327 B.15 GF20-C1 and GF20-C2: CO (1 − 0) spectra . . . 328 B.16 GF20-C1 and GF20-C2: 13_{CO (1 − 0) spectra . . . 329} B.17 GF20-C3 and GF20-C8: CO (1 − 0) spectra . . . 330 B.18 GF20-C3 and GF20-C8: 13_{CO (1 − 0) spectra . . . 331} B.19 Sa 187/188: CO (1 − 0) spectra . . . 333 B.20 Sa 187/188: 13_{CO (1 − 0) spectra . . . 334} B.21 Sa 187/188: C18_{O (2 − 1) spectra . . . 335} B.22 Sa 187/188: CS (2 − 1) spectra . . . 336 B.23 Sa 187/188: CS (3 − 2) spectra . . . 337 B.24 BHR 86: CO (1 − 0) spectra . . . 338 B.25 BHR 86: 13_{CO (1 − 0) and C}18_{O (2 − 1) spectra . . . 339} B.26 BHR 86: CS (2 − 1) and CS (3 − 2) spectra . . . 340 B.27 CB 63: CO (1 − 0) spectra . . . 342 B.28 CB 63: 13_{CO (1 − 0) spectra . . . 343} B.29 CB 63: C18_{O (2 − 1) spectra . . . 344} B.30 CB 63: CS (2 − 1) spectra . . . 345

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List of Tables

2.1 Observational parameters. . . 19 2.2 Nominal center positions. . . 20 2.3 Mean values of gaussian profile parameters. . . 23 2.4 Dense cores in GF 17, GF 20, and Sa 187/188. . . 66

3.1 Far-infrared properties of GF 17, GF 20, and Sa 187/188. . . 93

4.1 Basic properties of cylindrical and spherical polytropes. . . 125

5.1 _{CO (1 − 0) and} 13_{CO (1 − 0) peak line parameters. . . 138}

5.2 C18_{O J = 1 − 0 and J = 2 − 1 peak line parameters. . . 139}

5.3 _{CS J = 2 − 1 and J = 3 − 2 peak line parameters. . . 141} 5.4 Peak column densities derived from 13_{CO and C}18_{O observations. . . 151}

5.5 Sizes of the dense cores. . . 155 5.6 Results of velocity gradient fitting. . . 163 5.7 Energy balance of the dense cores. . . 183 5.8 Power law fits to the column density profiles. . . 196 5.9 Parameters of the best-fit polytropic sphere models. . . 201

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6.1 Results of gradient fitting. . . 210 6.2 Cloud masses and energy balance. . . 214 6.3 Isothermal cylinder best-fit parameters. . . 221 6.4 Parameters of the best-fit polytropic cylinder models. . . 226 6.5 Core separation. . . 240 6.6 Number of fragments formed in cylindrical clouds with L/D = 5. . . 242

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Chapter 1

Introduction

1.1 The Filamentary Nature of the Interstellar Medium

Understanding how stars form out of the diffuse interstellar medium (ISM) on both global and local scales is a fundamental open poblem in contemporary astrophysics (see McKee & Ostriker 2007 for a recent review). From the diffuse cirrus to the molecular clouds, onto the formation and death of stars, the galactic plane is the set where all the phases of the Galaxy life-cycle can be studied in context. More than 1000 dark clouds are visible within several hundred parsec of the Sun, and a substancial fraction of these contains young, low-mass stars (e.g., Lynds 1962; Herbig 1962). The great variety in cloud sizes, shapes, masses, complexity, and associated visible stars suggest that the observable dark clouds span a wide range of evolutionary stages, and that comparative studies of dark clouds and their stellar content may reveal important clues about cloud physics and the processes leading to star formation. Many of the nearest star-forming complexes have elongated structures of parsec scale, indicated by images at optical (Barnard 1927; Lynds 1962; Schneider & Elmegreen 1979; Dobashi et al. 2005), infrared (Lada, Alves, & Lombardi 2007), and submillimeter wave-lengths (Motte, Andr´e, & Neri 1998), and in rotational lines of CO and its isotopic variants (Ungerechts & Thaddeus 1987; Mizuno et al. 1995; Goldsmith et al. 2008). Several of these exhibit a “head-tail” structure (Tachihara et al. 2002), in which a central body of low aspect ratio and high column density is associated to a feature of greater aspect ratio and lower column density. With deeper observations and finer resolution, it now appears that some

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“heads” have more than one associated “tail”, and in particular, Myers (2009) has shown evidence that multiple parsec-scale filaments tend to branch out from dense star-forming “hubs” in regions forming stellar groups (e.g. L1495 in Taurus, NGC 7023, B59 in the Pipe Nebula, Serpens, L1688 in the Ophiuchus complex). Similar structures are also commonly seen at mid-infrared wavelengths with the Spitzer Space Telescope (Devine et al. 2004) in larger and more distant “infrared dark clouds” (IRDCs), now believed to be the precursors to massive stars and star clusters (Henning et al. 2010).

Very recently, the first results from the Gould Belt and Hi-GAL imaging surveys with the Herschel Space Observatory (Pillbrat et al. 2010) have revealed a profusion of parsec-scale filaments in galactic molecular clouds and suggest an intimate connection between the filamentary structure of the ISM and the formation process of dense cloud cores (Andr´e et al. 2010; Molinari et al. 2010). Remarkably, filaments are omnipresent even in unbound, non-starforming complexes such as the Polaris translucent cloud (Ward-Thompson et al. 2010; Miville-Deschˆenes et al. 2010). Furhermore, in active star-forming regions such as the Aquila Rift cloud, most of the prestellar cores identified with Herschel (Men’shchikov et al. 2010) are located within gravitationally unstable filaments for which the mass per unit length exceeds the critical value (Ostriker 1964).

Clearly, the physical processes responsible for these large-scale structures are acting over large distances, because the same processes appear at work in widely separated regions of the Galaxy. The prevalence of large-scale filamentary structures suggests that they may persist for a large fraction of a typical cloud lifetime, and therefore such structures may provide clues about the origin and geometry of star-forming regions.

1.2 The Origin of Filamentary Cloud Structures

The origin of filamentary cloud structure is unclear. Individual filaments can arise from compression of gas by converging turbulent flows (Jappsen et al. 2005; V´azquez-Semadeni et al. 2007; Heitsch et al. 2008), or by stellar wind bubbles, expanding H II regions, or supernova remnants (Whitworth 2007; Elmegreen 1998). Large scale maps of the HI gas in the Milky

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CHAPTER 1. INTRODUCTION 3

Way reveal a highly filamentary and shell-dominated structure that is probably produced by coherent supernova explosions and subsequent super-shell formation in the interstellar medium (Heiles 1997). It is possible that many elongated dark clouds condense out of dense HI filaments that form in the instabilities of the expanding surface as the swept-up gas, in the walls of these large, shocked and strongly magnetized shells (Heiles et al. 1993; McKee et al. 1993), cools down.

A detailed model of the origin of the Ophiuchus complex (de Geus 1992) suggests that massive stars in the Upper Scorpius and Upper-Centaurus Lupus subgroups originated wind-blown bubbles whose expansion swept up ambient gas to form the complex seen today (Mor-eira & Yun 2002 see also Chapter 6). In this model, OB winds combined with supernova explosions help to inflate the bubbles. Similarly, it has been suggested that both the Ophi-uchus and Lupus complexes may have formed as a consequence of the winds, ionization, and supernova shells from these subgroups of the Sco OB2 association (Tachihara et al. 2001). The molecular clouds of the Cyg X complex appear to have been shaped by winds and radia-tion from the OB stars in the Cyg OB2 associaradia-tion (Schneider et al. 2006). Similarly, it was suggested that the kinematics and structure of the CO gas in Orion A and B are consistent with stellar wind driven compression centered on Ori OB 1b (Wilson et al. 2005). Thus, several nearby star-forming regions appear to be close enough to OB subgroups to account for their compression by OB winds. Other sources of compression such as colliding turbulent flows and expanding H II regions may also play a role, but proximity to OB associations seems to have the greatest observational support.

1.3 Filamentary Dark Clouds and Star Formation

Filamentary Dark Clouds (FDC) or Globular Filaments (GF), are nearby (D < 900 pc) molecular clouds containing small numbers of dense cloud cores connected by low density gas and dust. These clouds look like strings (the low opacity material) with beads (the dense cores) strung along their lengths, often in a periodic fashion. Schneider & Elmegreen (1979), hereafter SE79, compiled a catalogue of 23 such globular filaments, and their intention was to establish an evolutionary link between ensembles of condensations, aligned along a

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filamentary dark cloud, and the smallest and simplest dark clouds in the Galaxy, the small, isolated Bok globules.

Bok globules are nearby small dark clouds of cold (Tkinetic ∼ 10 K) gas and dust often

containing small (∼ 0.1 pc), dense (nH2 ∼ 10

4 _cm−3_{) cores. Although Bok globules are}

relatively smaller than other dark clouds, in general their physical conditions seem to be very similar to those of dark clouds (Frerking & Langer 1982; Frerking, Langer, & Wilson 1987; Clemens & Barvainis 1988). In the classification proposed by Leung (1985), globules are divided into four types: elephant trunk and speck globules (Herbig 1974), cometary globules (Hawarden & Brand 1976), globular filaments (SE79), and isolated dark globules (Barnard 1919; Bok & Reilly 1947). Reipurth (1983) proposed that the different types of globules represent different evolutionary stages of these objects, beginning with cometary and elephant trunk globules and ending with isolated Bok globules. Indeed, this reflects the main motivation for the catalogue by SE79, who argue that globular filaments are the natural precursors of isolated globules. In this scenario, the FDCs most likely represent a transient phase of the molecular material, on its way from having resided in a larger cloud complex to becoming a collection of Bok globules, or dispersing altogether (Leung 1985).

Many filamentary dark clouds have dense cores engaged in the process of star formation. If self-gravity is important in FDCs, these dense cores are expected to be produced from a parent filamentary cloud by fragmentation and contraction, due to gravitational instabilities (Chandrasekhar & Fermi 1953). The core masses formed in this scenario are generally low (a few tens of a solar mass), and as such they are constrained to be potential sites for low mass star formation. The concerted formation of both cores and stars along the filamentary structures suggests a common trigger. If so, the dynamical and star formation history of globular filaments may indeed be a tracer of past events that are no longer directly observable. One may even think of several expanding supernova remnants disturbing the filamentary dark clouds, as pointed out by Launhardt (1996). Isolated Bok globules may trace such an event, too, if they are considered as part of a large-scale molecular cloud ensemble. Therefore, these filamentary structures have important and fascinating implications for the formation and evolution of galactic molecular clouds. However, more information about the

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three-CHAPTER 1. INTRODUCTION 5

dimensional nature of the filaments needs to be known before we can fully appreciate their origin and evolution.

Despite major and important progress in our knowledge of the star formation process in the Galaxy, we still lack a comprehensive theory of star formation. One reason why it is difficult to develop a convincing theory of star formation is that, so far, it has not been possible to define the initial conditions of a dense core that is to collapse into a star. In fact, these initial conditions cannot be defined until the formation of the isolated or regularly spaced dense cores is understood. A clear understanding of the initial conditions is necessary if we are to understand the processes by which clouds produce their star-forming cores. Thus, we believe that because they exhibit internal signs of fragmentation into a discrete collection of cores, filamentary dark clouds play an important role in the understanding of the initial stages of the star formation process.

The first step in this direction was undertaken by Chandrasekhar & Fermi (1953) and, since then, the theoretical framework was inspired by observations of isolated or globular filaments, and, reciprocally, motivated the search for evidence of fragmentation into conden-sations capable of forming stars. Several studies confirm the ubiquity of these filaments in molecular clouds such as Taurus (Gomez & Pudritz 1992), and the R -shaped filament in Orion (Dutrey et al. 1991), and have revealed that many molecular clouds are filamentary structures with nearly periodic density enhancements that are supported by non-thermal, small-scale MHD motions of some kind, as well as large scale ordered magnetic fields (cf. Schleuning 1998).

The problem of fragmentation of self-gravitating filaments has accordingly received some attention, and since the early work of Chandrasekhar & Fermi (1953) several theoretical approaches were undertaken on filamentary structures. The linear stability of a cylindrical cloud of gas has been investigated extensively by Stod´olkiewicz (1963), Ostriker (1964, 1965), Larson (1985), Nagasawa (1987), Matsumoto, Nakamura, & Hanawa (1994), Nakamura, Hanawa, & Nakano (1993, 1995), Nakajima & Hanawa (1996), Fiege & Pudritz (2000a,b,c); Boss (2007) and Oproiu & Horedt (2008). These studies take into account the effects of mag-netic fields, stiffness of the equation of state, rotation, and collapse of the filamentary cloud

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in the radial direction. The fragmentation of isothermal cylinders has been also extensively studied by numerical simulations (Jappsen et al. 2005, and references therein). Bastien (1983) found that the fragmentation of non-rotating, non-magnetic isothermal cylindrical clouds is marginal when the aspect ratio L/R is ∼ 4, but very marked when L/R ∼ 10. Stenholm & Pudritz (1993) started from a uniform, magnetically subcritical medium, agi-tated by a spectrum of Alfv´en waves. The cores forming in the course of time, as resulting from a N-body MHD code, exhibit sizes which compare well with the observations (a few tenths of a parsec). Monaghan (1994) simulated the role of vorticity and angular momentum on the fragmentation of an initially homogeneous, non-magnetized ellipsoid, and obtains a sequence of regularly spaced globules, with a spacing comparable to the wavelength of ma-ximum growth predicted by Larson (1985). Contrary to these results, the simulations of Bastien et al. (1991), also neglecting magnetic fields, show binary structure at the ends of the initially cylindrical filament, filled with interstitial clumps. Nakamura, Hanawa, & Nakano (1993, 1995) and Fiege & Pudritz (2000a,b) resumed the problem of the stability of isothermal magnetized filamentary clouds. Due to instabilities, a gas flow induced along the magnetic field lines leads to the formation of disks, that subsequently collapse to pro-tostellar condensations (see also the “pseudo” disks of Galli & Shu 1993a,b). These papers consitute the most exhaustive examination of the stability of magnetized, compressible, self-gravitating filaments with realistic density profiles (i.e., density profiles much shallower than the r−4 _{solution of Ostriker 1964).}

Despite the large number of observations of FDCs, complete observational studies of FDCs at infrared and millimeter wavelengths probing their physical properties are still lacking. With the advent of large-scale imaging surveys at near infrared wavelengths such as the Two Micron All Sky Survey (2MASS)1 _{it became possible to derive dust extinction maps of}

star-forming complexes (Lombardi, Alves & Lada 2006; Lombardi, Lada & Alves 2008b) at a few (∼ 3) arcmin of spatial resolution. On the other hand, most of the millimeter observations performed on nearby FDCs are part of large-scale CO surveys of dark clouds (Bally et al. 1987; Dobashi et al. 1994; Tachihara et al. 1996; Mizuno et al. 1998; Kawamura et al. 1998)

1_{The Two Micron All Sky Survey project is a collaboration between the University of Massachusetts and the Infrared}

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and dense cores (Vilas-Boas, Myers, & Fuller 1994; Onishi et al. 1998; Muench et al. 2007; Foster et al. 2009). As such, most of these observations were carried out with poor spatial resolution due to the large beamsizes used (typically greater than ∼ 3 arcmin). Furthermore, the existing pointed observations toward FDCs very often do not completely cover the full extent of the clouds (e.g. GF 9, Ciardi et al. 2000) simply because observations of large areas of the sky with spatial resolutions < 1 arcmin are extremely time consuming. Thus, these molecular line studies often suffer from low spatial resolution, incomplete sampling, and, due to the limitations of the current instrumentation, they are typically available only for small regions of the sky. This situation, however, has been rapidly changing with the development of focal plane array instrumentation.

Because many star-forming molecular clouds are filamentary, most of the work performed on FDCs has concerned the star formation properties of individual dense cores (see Bergin & Tafalla 2007 for a review) within FDCs, and in many cases, the studies have been part of general surveys of low-mass star forming regions and have not represented investigations of an individual FDC (e.g., Myers & Benson 1983; Myers, Linke, & Benson 1983; Beichman et al. 1986; Wiesemeyer 1997; Foster et al. 2009). The “dense cores” within dark clouds represent a relatively narrowly defined and homogeneous subset of dark cloud gas, and the importance of such cores in forming low-mass stars seems well established. However, in order to investigate how dark clouds produce both cores and stars, it is necessary to study a much wider variety of dark cloud regions than cores alone. Therefore, many individual low-mass star-forming cores contained within FDCs have been studied, but the FDCs, as entities of their own, are still poorly known.

1.4 Aim of the Thesis

Given this panorama, the need for a comparative study of physical conditions within FDCs became evident, as detailed observations of filamentary dark clouds will likely bring new insight into the nature of these objects. By conducting such a study on a sample of FDCs, one can hope to learn about the formation, evolution, and destruction of filamentary dark clouds by studying the structural and kinematical similarities and differences between FDCs.

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This may provide important clues about their origins, but also significant constraints on the existing poorly constrained theoretical models of filamentary cloud stability and collapse.

Gas density is one of the fundamental properties of a molecular cloud, since it is related to the cloud’s gravitational potential, cooling rate, and evolutionary state, and since it determines the basic time and length scales for star formation in the cloud. Thus, in order to understand the evolution of filamentary dark clouds one needs to understand their density structure, which implies assessing the large-scale distribution of gas density in the clouds. This has the advantage of simultaneously yielding an unbiased survey of dense cores and their spatial distribution within these clouds, as a side product. Since stars form in dense cores, one might expect the spatial distribution of cores to reflect the distribution of star formation within a molecular cloud and, most important, to locate possible sites of recent or future star formation. Furthermore, it also provides a census of core properties whithin a cloud, and this information can be used to determine the size, shape, and mass spectra of the cores yielding valuable insights toward the understanding of the processes, such as fragmentation, that govern the formation and evolution of these clouds.

We thus decided to conduct such an investigation in a small sample of three filamentary dark clouds. The observational data elements of our study were (i) millimeter molecular line maps of the FDCs in the rotational transitions of CO, 13CO, C18O, and CS; and (ii) IRAS2 Sky Survey Atlas (ISSA) co-added images of the three FDCs at 12, 25, 60, and 100 µm. Additional millimeter observations in the same molecular transitions were obtained for two Bok globules: these add to the available data from the literature on surveys of isolated Bok globules, which we can use in order to explore a comparison between the dense cores within our filamentary clouds and the small Bok globules.

Observations of the CO and CS line transitions allow direct determination of the gas kinematics, density, and temperature, and the far-infrared images allow us to analyse the dust content of each cloud. Using these observations, we address a number of questions, such as the following:

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• Are the filamentary dark clouds stable? Is there enough mass to gravitationally bind the gas motions present? If not, are these clouds confined by an external pressure? • Is rotation important in these filaments? If so, how does it affect continuing

fragmen-tation?

• What are the radial gas density profiles of filaments? What are their masses per unit length?

• What fraction of the mass of each filament is contained in dense cores?

• How does the distribution of the dust correlate with the distribution of the molecular gas within each cloud?

• What is the nature of the dense cores in FDCs? Are their physical properties similar to the ones found for other dark cloud cores, or to those of Bok globules?

• What are the radial gas temperature and density distributions in the dense cores re-gions? Are these regions relatively quiescent, with cold cores, as is generally seen in starless Bok globules (Dickman & Clemens 1983)? Are the gas density profiles as steep as for the Bok globules, indicative of fairly old ages, or are they shallower, indicating that the cores are perhaps still accumulating gas and dust?

• What is the nature of the filamentary material between the dense cores?

• Are the gas motions in the cores thermal or do they present a non-negligible contribution from turbulence (non-thermal line-broadening, e.g. Myers & Fuller 1992; Foster et al. 2009)?

The questions raised above reveal the central aim of the present thesis, which can be broken in a four-fold manner: (i) assess the structure and kinematics of the filamentary dark clouds, (ii) investigate their stability by determining and modelling their density profiles, (iii) analyse their dense gas content, and establish their dense core properties and star formation activity, and (iv) establish the source of dust heating in these clouds.

Altogether, this study allows us to investigate the origins of these clouds, to obtain density profiles on the size scales of real filamentary clouds, not just of the individual dense cores

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within, and finally to compare the dense core properties with those of isolated Bok globules. In short, the present work represents an effort to bring together the theory and observations of filamentary dark clouds.

1.5 Overview of the Subject Clouds

1.5.1 Selection Criteria

Our sample consists of the three filamentary dark clouds GF 17 (also known as Lupus 4), GF 20 (also known as Lupus 2) and Sa 187/188 (also know as the Norma cloud). GF 17 and GF 20 were selected from the catalog of 23 globular filaments of SE79, and Sa 187/188 was selected from the catalog of Sandqvist (1977). The criteria used to select our sources were six-fold:

• the filamentary dark clouds were required to have modest long axes (≤ 1 degree) and resolvable short axes (at least several arcminutes).

• IRSKY3 _{was used to select a sub-sample of FDCs for which embedded intermediate}

mass star formation was not active (since that tends to alter the local cloud environment drastically). The resulting sample contains some low mass star formation, as indicated by the presence of some IRAS Point Source Catalog (PSC) sources.

• the sample should contain filamentary clouds in different stages of the star formation activity, i.e., should include at least one cloud which is relatively quiescent and one cloud with signposts of recent or present star formation.

• the final sample should include nearby filamentary dark clouds, at different distances from the Sun, in order for different linear sizes to be resolved.

• filamentary dark clouds which were poorly studied in the past were preferentially se-lected. We have used the SIMBAD4 _{database to search for previous existing studies of}

these sources (see section 1.5).

3_{IRSKY is an observation planning tool for the infrared sky developped at the Infrared Processing and Analysis Center}

(IPAC), NASA’s multi-mission center of expertise for long-wavelength astrophysics.

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We have also observed Bok globules CB 63 (Clemens & Barvainis 1988) and BHR 86 (Bourke, Hyland, & Robinson 1995). While CB 63 is a quiescent starless globule (Clemens, Dickman, & Ciardi 1992), BHR 86 is known to be actively engaged in star formation, as indicated by the presence of a protostar, possibly a Class 0 (Lehtinen 1997).

1.5.2 GF 17 and GF 20

The Lupus molecular cloud complex is one of the nearest sites of low-mass star formation. It is located in the region with galactic coordinates 334◦

< l < 352◦

and +5◦

< b < +25◦

, and consists of 9 loosely connected subgroups designated as Lupus 1 through 9. For reference, Lupus 1 is B228, while Lupus 2 is GF 20, and Lupus 4 is GF 17. These clouds are located near the Sco OB2 association (Humphreys 1978), and there is no indication that massive OB stars are formed. A detailed view of the complex can be appreciated in Figure 10 of Lombardi, Lada & Alves (2008b) obtained from near infrared dust extinction maps based on 2MASS data.

Lupus is a well studied complex showing different modes of star formation: Lupus 1 shows isolated star formation, while cluster formation is seen in Lupus 3 (Schwartz 1977; Hughes et al. 1994; Nakajima et al. 2000), and no evident star formation is seen in Lupus 5. Most of the optically visible young stars in Lupus were first identified on the basis of their Hα emission in a deep objective prism survey by Schwartz (1977). Since then, the YSO population of the Lupus star-forming region has been assessed: spectroscopic and infrared photometric observations of the T Tauri population were carried out by Hughes et al. (1994). Krautter et al. (1997) and Wichmann et al. (1997) studied 132 weak-line T Tauri stars discovered in ROSAT5 _{images of Lupus. Nuernberger, Chini, & Zinnecker (1997) performed a 1.3 mm}

dust continuum survey of Hα selected T Tauri stars. The age distribution of the T Tauri stars suggests that this star forming region is relatively old, ∼ 107 _{yr (Hughes et al. 1994).}

Six Herbig-Haro objects are found in Lupus (Krautter 1991; Reipurth 1994), while a single molecular outflow, a sign of recent star formation, is known (Tachihara et al. 1996).

The Lupus cloud complex was observed as part of the Spitzer Legacy Project “From

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Molecular Cores to Planet-forming Disks” (c2d; Evans et al. 2003) and, recently, Mer´ın et al. (2008) reported on such observations of Lupus 1, 3 and 4 with the Spitzer Space Telescope Infrared Array Camera (IRAC) and the Multiband Imaging Photometer (MIPS). Lupus 2, however, was not observed by c2d. They find 4 confirmed Class II YSOs, and suggest the existence of an additional 7 Class II YSO candidates and a Class I YSO candidate, all seen toward the “main core” region of GF 17 (see their Figure 22). More recently, Comer´on, Spezzi & Mart´ı (2009) discovered a new population of 130 cool members (both stars and brown dwarfs that have lost their inner disks on a timescale of a few Myr or less) in regions of low extinction of Lupus 1 and 3, and noted that this population is absent in Lupus 4. On the other hand, only five T Tauri stars are known to be associated with GF 20, one of these being the extremely active T Tauri star RU Lupi, which is seen toward the western condensation, about 30 _{to the NE of the Herbig Haro object HH 55.}

Figure 1.1 shows Digital Sky Survey (DSS) images of the filamentary dark clouds GF 17 and GF 20. The latter cloud appears as a small filament in which one can identify several condensations. In turn, GF 17 is a much larger filament, also with several identifiable condensations. Due to its large area on the sky (∼ 20◦

), millimeter observations have mostly covered very limited regions of the Lupus complex. The first extensive cloud survey in the CO (1 − 0) line over an area of ∼ 170 square degrees was made by Murphy, Cohen, & May (1986) with an effective resolution of 300

. They estimated the mass of the cloud complex to be 3 × 104 _M

, comparable to that of the nearby Ophiuchus clouds. Gahm, Johansson,

& Liseau (1993) made CO (1 − 0) observations of GF 20 with 4000

resolution but covering only the main condensation of the cloud. Using optical star counts, Andreazza & Vilas-Boas (1996) derived masses of 197 M and 59 M for GF 17 and GF 20, respectively. More

recently, Hara et al. (1999) and Tachihara et al. (2001) have performed a survey in CO (1 − 0) and C18_{O (1 − 0) using the NANTEN sub-millimeter telescope. Their data covered}

an area of 550 square degrees with an effective resolution of 80

, and smaller areas with 2.30

resolution. The Lupus clouds were also searched for the presence of magnetic fields (Myers & Goodman 1991).

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150 ± 20 pc (Krautter 1991), has been challenged by Knude & Hog (1998), who reported 100 pc, and by Wichmann et al. (1998), 190 pc, both using Hipparcos observations. In addition, Knude & Nielsen (2001) suggested that Lupus 2 (GF 20) is physically disconnected from the other Lupus subclouds, and is located at a distance of 360 pc. However, surprisingly no velocity difference is found among the various Lupus subclouds in radio observations. More recently, Lombardi, Lada & Alves (2008a) used a maximum-likelihood technique based on Hipparcos data obtaining d = 155 ± 8 pc. For consistency with the previous work by Moreira & Yun (2002), we will adopt the value of 150 pc as the actual distance of Lupus 2 and 4, in agreement with Mer´ın et al. (2008).

1.5.3 Sa 187/188

The southern constellation Norma contains a number of small dark clouds. One of these, Sa 187/188, located around α ≈ 16h_{30 δ ≈ −45}o _{(Sandqvist 1977) and known as the Norma}

cloud, is a small elongated dark cloud consisting of two main parts, a larger, structured “head” to the west (Sa 187) and a narrow (about 40

in diameter) filamentary “tail” (Sa 188) stretching to the east (Figure 1.2), with the east-west extent of the whole cloud being about 400

. There is no direct determination of the distance of the cloud: Cohen et al. (1984) suggested 200 pc, Reipurth (1985) favored 300 pc, Graham & Frogel (1985) argued for 700 pc, Cohen, Dopita, & Schwartz (1986a) used 940 pc, and Gregorio-Hetem, Sanzovo, & L´epine (1988) proposed 150 pc. Re-examining the various arguments, a distance of 500−700 pc seems now most probable. For the purpose of the present dissertation we will adopt the distance of 600 pc, noting with these authors the considerable uncertainty in this estimate. Four IRAS point sources have been found inside the Sa 187/188 complex.

Sa 187 is known to be a site of active star formation, and has attracted considerable interest since Schwartz (1977) discovered two Herbig-Haro objects, HH 56 and HH 57, within the main core region. HH 57 was found to be driven by the near-infrared source IRS 8 (Reipurth & Wamsteker 1983), while HH 56 appears to be driven by a young star embedded in the small reflection nebula Re 13 (Alvarez et al. 1986). Graham & Frogel (1985) found that, within this same region, the infrared source V346Nor (first detected by Elias 1980) is a

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Figure 1.2: Digital Sky Survey image of the filamentary dark cloud Sa 187/188.

FU Ori type star (Prusti et al. 1993; Reipurth et al. 1997; ´Abrah´am et al. 2004). Moreira et al. (2000) found a second star-forming core towards the IRAS point source IRAS 16295−4452 (see Chapter 7).

Sa 188 (the eastern filament) is also presently forming stars; a nebulous star is located in the middle of the filament, two small reflection nebulae are found in each end of the filament (Reipurth et al. 1993), and two young IRAS sources associated with small molecular outflows were identified by Reipurth et al. (1997). Thus, several young stars in Sa 187/188 are known in different stages of evolution, with the Sa 188 being less actively engaged in star formation than Sa 187.

The large-scale extinction patterns in this part of the Milky Way have been discussed by Gregorio-Hetem, Sanzovo, & L´epine (1988). CO observations were made of the cloud complex by Alvarez et al. (1986), which suggest a total mass of 500 (D/700)2 _M

. More

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of Sa 187/188, taken with the SEST6 _{Imaging Bolometer Array (SIMBA). They have found}

that the distribution of the dust extended millimeter continuum emission follows very well the absorption structure visible on DSS images, and derived a total H2 mass of 340 M,

assuming a distance of 700 pc. In adition, they have found six compact millimeter sources, five of which coincide with known stellar objects.

1.6 Thesis Plan

The remainder of this dissertation is presented in seven chapters. Chapter 2 presents the CO and CS millimeter spectroscopic survey observations and the molecular line maps obtained for GF 17, GF 20, and Sa 187/188; it describes our LTE method of analysis of the molecular line data, and finally presents the dense cores identified in this survey. In Chapter 3 we present IRAS co-added images and discuss the correlations between gas and dust in GF 17, GF 20, and Sa 187/188. In particular, we show that GF 17 and GF 20 are externally heated clouds, unlike Sa 187/188. Chapter 4 describes the models of spherical and cylindrical polytropic clouds used to model the observed column density profiles of both the dense cores, and the filamentary structures as a whole. In Chapter 5 we characterize the properties of the dense cores in GF 17, GF 20, and Sa 187/188, and establish comparisons with those of Bok globules. The kinematics, energetics, and structure of our clouds are discussed in Chapter 6, and Chapter 7 reports the discovery of new evidence for previously unknown star formation activity in GF 20 and Sa 187/188. Finally, Chapter 8 closes the document by summarizing the work involved in this dissertation and presenting suggestions to explore in future research.

6_{The Swedish-ESO Submillimeter Telescope, de-commissioned in August 2003, was operated jointly by ESO and the Swedish}

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Chapter 2

Millimeter Molecular Line Mapping

2.1 Introduction

The bulk of the mass in molecular clouds is in the form of molecular hydrogen (H2) (Genzel

1992). Unfortunately, two factors render H2 largely unobservable in the cold dense

environ-ments of molecular clouds (e.g., Lada et al. 1994): (1) being an homonuclear molecule, H2

lacks a permanent dipole moment and has extremely weak rotational transitions, and (2) being the lightest interstellar molecule, its lower energy rotational transitions are at mid-infrared wavelengths which are both inaccessible to observations from earth and also too energetic to be collisionally excited at the low temperatures (T ≈ 10 K) that characterize these clouds. Hence, spectroscopic observations of other molecules which are environmen-tally coincident with H2 (i.e., molecules which trace molecular hydrogen) are required to

assess the gas properties inside molecular clouds (Turner 1988).

Excluding H2, which totally lacks a radio spectrum, carbon monoxide (12C16O, hereafter

CO) is the most abundant molecular constituent of dark clouds. Since its first astronomi-cal detection by Wilson, Jefferts, & Penzias (1970), the widespread and easily observable λ2.6mm J = 1 − 0 rotational transition of CO has become an extremely valuable tool in the study of the more highly condensed parts of the interstellar medium, i.e. molecular clouds. Indeed, observations of this molecule have proven useful, for example in studies of the gas kinetic temperature, cloud morphology, and search for molecular outflows within star-forming regions. While CO traces gas with molecular hydrogen densities of ∼ 102 _cm−3_,