THE HISTORY OF CHEMICAL ENRICHMENT AND THE SITES OF EARLY
NUCLEOSYNTHESIS: CNO ABUNDANCES OF GALACTIC
CARBON-ENHANCED METAL-POOR STARS
By
Catherine R. Kennedy

A DISSERTATION
Submitted to
Michigan State University
in partial fulfillment of the requirements
for the degree of
DOCTOR OF PHILOSOPHY
Physics and Astronomy
2011

ABSTRACT
THE HISTORY OF CHEMICAL ENRICHMENT AND THE SITES OF EARLY
NUCLEOSYNTHESIS: CNO ABUNDANCES OF GALACTIC
CARBON-ENHANCED METAL-POOR STARS
By
Catherine R. Kennedy

This dissertation focuses on abundance analyses of carbon-enhanced metal-poor (CEMP)
Galactic halo stars. Different methods for determining carbon, nitrogen, oxygen, and
also some barium abundances are described. The study of these abundances in such
stars serves to investigate the means by which the Universe became enriched in metals. Due to the different kinds of CEMP stars observed in the Milky Way, it can only
be assumed that there is certainly more than one method of carbon-enhancement at
early times. Complete abundance analyses for as many of these archaeological relics
as possible are needed in order to constrain the astrophysical sites of early carbon
production.
There are three main parts of this dissertation. The first part describes new techniques to determine oxygen abundances from spectra of the near-infrared molecular
CO bands. With the near-IR OSIRIS spectrograph on the SOAR 4.1-m telescope, 57
CEMP stars were observed. A wide range of oxygen abundances were estimated, and
the results were statistically compared to high-resolution estimates for both carbonenhanced and carbon-normal metal-poor stars. Abundance patterns of the sample
stars were compared to yield predictions for very metal-poor asymptotic giant branch
(AGB) stars. The majority of the sample exhibit patterns consistent with CEMP
stars having s-process-element enhancements, and thus have very likely been polluted by carbon- and oxygen-enhanced material transferred from a metal-poor AGB

companion.
The second part delineates a new survey effort implemented in order to identify
new CEMP stars. For the initial pilot study, a new selection technique was developed
based solely on the strength of the CH G band at 4300 ˚. This technique eliminated
A
previous temperature and metallicity biases present in other CEMP surveys. Observations of the pilot sample were carried out with the Goodman HTS spectrograph
on the SOAR 4.1-m telescope. Of the over 120 candidate stars observed, over 35%
were found to be CEMP stars. The selection technique was then improved to include
a second index for the strength of the G band, and the survey was continued on
both the SOAR and Gemini telescopes. After this extension, the success rate of this
program increased to 50%.
The final part of this dissertation contains details of a pilot study of known metalpoor stars using the X-Shooter spectrograph on the Very Large Telescope (VLT).
With three spectrograph arms (near-ultraviolet, optical, and near-infrared), this instrument was used to calculate carbon, nitrogen, oxygen, and barium abundances for
a sample of 27 CEMP stars. The broad spectral range of this instrument is unprecedented, and it is an efficient way to estimate abundances for several pertinent species
in just one exposure per star. Of the 27 stars, many proved to be enhanced in carbon.
The majority appear to be consistent with s-process-element enhancement, but there
was one extremely metal-poor star which falls into the rare family of CEMP stars
with no neutron-capture-element enhancement.

Copyright by
CATHERINE R. KENNEDY
2011

To my mother and father for the love and support.
To Lauren, Nora, Liz, Amanda, Jennie, Cait, Val, Bill, and Dave for the
lasting friendship and the laughter.
To the person I have looked up to thoughout my entire life, John. I have always
endeavored to be his equal with regard to integrity, work ethic, and academic
success. He is as fast now as he was on his Big Wheel; but if I pedal fast, I may
catch up to him.

v

ACKNOWLEDGMENTS

Many thanks to Timothy C. Beers for his guidance for the past five years. Special
thanks to Silvia Rossi, Vinicius Placco, Thirupathi Sivarani, and Young Sun Lee for
endless amounts of collaborative spirit and scientific expertise. C.R.K. acknowledges
partial support for this work from grants AST 07-07776, as well as from PHY 02-15783
and PHY 08-22648; Physics Frontier Center/Joint Institute for Nuclear Astrophysics
(JINA), awarded by the US National Science Foundation. Thanks to Christopher
A
Waters for the L TEX class used to format this thesis.

vi

TABLE OF CONTENTS

List of Tables. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . ix
List of Figures. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . xi
List of Symbols . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . xv
1 Introduction. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.1
The Milky Way Components and Formation Scenarios . . . . .
1.2
The Carbon-Enhanced Metal-Poor Stars of the Galactic Halo .

1
2
3

2 Oxygen Abundances for CEMP Stars from Near-Infrared Spectroscopy
with SOAR/OSIRIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7
2.1
Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7
2.2
Observations and Data Reduction . . . . . . . . . . . . . . . . 10
2.3
Adopted Atmospheric Parameters and Synthetic Spectra . . . . 11
2.4
Determination of [O/H] . . . . . . . . . . . . . . . . . . . . . . 15
2.5
Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17
2.5.1
Statistical Comparison to High-resolution Oxygen Estimates 20
2.5.2
High-resolution Nitrogen Estimates . . . . . . . . . . . . . . 21
2.6
Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24
2.6.1
[O/Fe] in Carbon-normal and Carbon-enhanced Metal-poor
Stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 25
2.6.2
C, N, and O: Comparison with AGB Models . . . . . . . . . 26
2.6.3
Considering the Effects of Dilution . . . . . . . . . . . . . . 29
2.6.4
Uncertainties of the AGB Models . . . . . . . . . . . . . . . 31
2.7
Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 32
3 A Survey to Identify New CEMP Stars . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.1
Motivation . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2
The Selection Technique: A New Extended Index for the CH G
Band . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3
Target Selection . . . . . . . . . . . . . . . . . . . . . . . . . .
3.4
Observations and Data Reduction . . . . . . . . . . . . . . . .
3.5
Atmospheric Parameters and Carbon Abundances . . . . . . .
3.6
Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.7
Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
vii

34
34
35
36
40
47
54
55

4 Improving the Selection Technique and Extending the Survey with
Multiple Observatories . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1
Refining the Selection Criteria . . . . . . . . . . . . . . . . . .
4.2
New Selection . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3
Follow-up Observations: GMOS on Gemini . . . . . . . . . . .
4.4
Follow-up Observations: Goodman HTS on SOAR . . . . . . .

57
57
59
61
66

5 Carbon, Nitrogen, Oxygen, and Barium Abundances with XSHOOTER
on VLT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71
5.1
Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71
5.2
Selected Targets and Observations . . . . . . . . . . . . . . . . 72
5.3
Model Atmospheres, Line Lists, and Synthetic Spectra . . . . . 73
5.3.1
Model Atmospheres . . . . . . . . . . . . . . . . . . . . . . 73
5.3.2
Line Lists . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73
5.3.3
Synthetic Spectra . . . . . . . . . . . . . . . . . . . . . . . . 74
5.4
Abundance Determinations . . . . . . . . . . . . . . . . . . . . 74
5.5
Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76
5.6
Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 79
5.6.1
Neutron-capture-enhanced Stars . . . . . . . . . . . . . . . 79
5.6.2
CEMP-no Stars . . . . . . . . . . . . . . . . . . . . . . . . . 82
5.7
Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83
6 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84
A XSHOOTER Spectra: [C/Fe], [N/Fe], [O/Fe], and [Ba/Fe] Determinations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 88
A.1
[C/Fe] . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 88
A.2
[N/Fe] . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102
A.3
[O/Fe] . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 116
A.4
[Ba/Fe] . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 125
B Acronyms . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 133
References . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 135

viii

LIST OF TABLES

1.1

Classification of CEMP Stars . . . . . . . . . . . . . . . . . . . . . . . .

4

2.1

Atmospheric Parameters and Carbon Abundances I . . . . . . . . . . . .

13

2.2

Atmospheric Parameters and Carbon Abundances II . . . . . . . . . . .

14

2.3

Atmospheric Parameters and Carbon Abundances III . . . . . . . . . . .

15

2.4

Abundance Ratios for the Entire Sample I . . . . . . . . . . . . . . . . .

22

2.5

Abundance Ratios for the Entire Sample II . . . . . . . . . . . . . . . . .

23

2.6

Abundance Ratios for the Entire Sample III . . . . . . . . . . . . . . . .

24

3.1

Classification of CEMP Candidates . . . . . . . . . . . . . . . . . . . . .

38

3.2

Color, GPE, and Category of Pilot Stars I . . . . . . . . . . . . . . . . .

41

3.3

Color, GPE, and Category of Pilot Stars II . . . . . . . . . . . . . . . . .

42

3.4

Color, GPE, and Category of Pilot Stars III . . . . . . . . . . . . . . . .

43

3.5

Color, GPE, and Category of Pilot Stars IV . . . . . . . . . . . . . . . .

44

3.6

Color, GPE, and Category of Pilot Stars V . . . . . . . . . . . . . . . . .

45

3.7

Color, GPE, and Category of Pilot Stars VI . . . . . . . . . . . . . . . .

46

3.8

Atmospheric Parameters and Carbon Abundances for Pilot Stars I . . . .

49

3.9

Atmospheric Parameters and Carbon Abundances for Pilot Stars II . . .

50

3.10 Atmospheric Parameters and Carbon Abundances for Pilot Stars III . . .

51

3.11 Atmospheric Parameters and Carbon Abundances for Pilot Stars IV . . .

52

3.12 Atmospheric Parameters and Carbon Abundances for Pilot Stars V . . .

53

3.13 Atmospheric Parameters and Carbon Abundances for Pilot Stars VI . . .

54

4.1

63

Atmospheric Parameters and Carbon Abundances for Gemini Stars I . .
ix

4.2

Atmospheric Parameters and Carbon Abundances for Gemini Stars II . .

64

4.3

Atmospheric Parameters and Carbon Abundances for Gemini Stars III .

65

4.4

Atmospheric Parameters and Carbon Abundances for Goodman Stars I .

67

4.5

Atmospheric Parameters and Carbon Abundances for Goodman Stars II

68

4.6

Atmospheric Parameters and Carbon Abundances for Goodman Stars III

69

5.1

Atmospheric Parameters and Barium Abundances for X-Shooter Stars I .

76

5.2

Atmospheric Parameters and Barium Abundances for X-Shooter Stars II

77

5.3

Nitrogen, Carbon, and Oxygen Abundances for X-Shooter Stars I . . . .

77

5.4

Nitrogen, Carbon, and Oxygen Abundances for X-Shooter Stars II . . . .

78

x

LIST OF FIGURES

1.1

Path of s- and r-process . . . . . . . . . . . . . . . . . . . . . . . . . . .

5

2.1

Three synthetic spectra with different [O/Fe] ratios. . . . . . . . . . . . .

17

2.2

Estimates of [O/Fe] for four stars. . . . . . . . . . . . . . . . . . . . . . .

18

2.3

Distribution of [O/Fe] . . . . . . . . . . . . . . . . . . . . . . . . . . . .

19

2.4

[O/Fe] vs. [Fe/H] for stars with [C/Fe] ≥ +1.75. . . . . . . . . . . . . . .

20

2.5

[O/Fe] compared to high-resolution analysis . . . . . . . . . . . . . . . .

21

2.6

C, N, and O abundances vs. metallicity for 10 stars . . . . . . . . . . . .

27

2.7

Abundances compared to AGB abundance yields . . . . . . . . . . . . .

28

2.8

The effects of dilution on abundances . . . . . . . . . . . . . . . . . . . .

30

3.1

The new GPE index . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

36

3.2

GPE vs. (J−K)0 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

37

3.3

Index-color distribution of CEMP candidates of different types . . . . . .

39

3.4

Metallicity cutoff for CEMP candidates . . . . . . . . . . . . . . . . . . .

40

3.5

Goodman spectra of pilot-program stars I . . . . . . . . . . . . . . . . .

46

3.6

Goodman spectra of pilot-program stars II . . . . . . . . . . . . . . . . .

47

3.7

Goodman spectra of pilot-program stars III . . . . . . . . . . . . . . . .

48

3.8

[C/Fe] estimate for pilot-program star. . . . . . . . . . . . . . . . . . . .

49

3.9

[C/Fe] versus [Fe/H] for Pilot Sample. . . . . . . . . . . . . . . . . . . .

55

4.1

The new EGP index for 6 HES stars. . . . . . . . . . . . . . . . . . . . .

58

4.2

Distribution of GPE and EGP indices for both the stars and the bright
subsets. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

60

xi

4.3

Distribution of GPE and EGP indices after the saturation correction. . .

61

4.4

Selection criteria for GPE and EGP indices. . . . . . . . . . . . . . . . .

62

4.5

[C/Fe] versus [Fe/H] for Pilot Sample + Gemini Sample. . . . . . . . . .

66

4.6

[C/Fe] versus [Fe/H] for Pilot + Gemini + Additional Goodman samples.

70

5.1

C, N, O, and Ba spectral synthesis. . . . . . . . . . . . . . . . . . . . . .

75

5.2

C, N, O, and Ba abundances. . . . . . . . . . . . . . . . . . . . . . . . .

79

5.3

Abundance patterns for CEMP X-Shooter stars. . . . . . . . . . . . . . .

80

5.4

CNO Abundances of metal-poor AGB Stars . . . . . . . . . . . . . . . .

81

A.1 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

88

A.2 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

89

A.3 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

90

A.4 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

91

A.5 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

92

A.6 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

93

A.7 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

94

A.8 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

95

A.9 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

96

A.10 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

97

A.11 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

98

A.12 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . .

99

A.13 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 100
A.14 Carbon abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 101
A.15 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 102
A.16 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 103
xii

A.17 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 104
A.18 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 105
A.19 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 106
A.20 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 107
A.21 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 108
A.22 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 109
A.23 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 110
A.24 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 111
A.25 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 112
A.26 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 113
A.27 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 114
A.28 Nitrogen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 115
A.29 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 116
A.30 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 117
A.31 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 118
A.32 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 119
A.33 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 120
A.34 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 121
A.35 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 122
A.36 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 123
A.37 Oxygen abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . 124
A.38 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 125
A.39 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 126
A.40 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 127
A.41 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 128
xiii

A.42 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 129
A.43 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 130
A.44 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 131
A.45 Barium abundances for X-Shooter stars . . . . . . . . . . . . . . . . . . . 132

xiv

LIST OF SYMBOLS

unit of distance equal to 3.0857 × 1018 cm . . . . . . . . . . . . . . . . .

2

Teff effective temperature . . . . . . . . . . . . . . . . . . . . . . . . . . . .

11

g

surface gravity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

11

M⊙ solar mass . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

26

pc

xv

Chapter 1:
Introduction
There are a variety of methods by which the Universe was polluted with all of the
elements that we observe today. Big Bang nucleosynthesis provided the early Universe with only hydrogen, helium, and trace amounts of lithium. All other elements
have since been created in stars. As early generations of stars exploded as supernovae, the interstellar medium (ISM) was polluted with the heavy elements that were
created within them and during their explosions. Subsequent generations of stars
were then born of this more chemically-enriched material. This cycle of the birth and
death of generations of stars continues today, and the main result of this process is
that the generations of stars become more and more chemically-enriched than their
predecessors as time goes on.
In order to quantify this process, one can measure the amount of metals in stars
via stellar spectroscopy. In astronomical terms, “metals” are defined as all elements
except for hydrogen and helium. In this dissertation, the metallicity of a star is defined
as [Fe/H], a logarithmic comparison of the amount of metal in a star compared to
that of the Sun. Other abundances will be discussed in terms of the same notation:

[A/B] = log(NA /NB )star − log(NA /NB )⊙ ,

(1.1)

where NA and NB are the number densities of the elements A and B for the stars.
In these terms, a star which has [Fe/H] = −1.0 has a metallicity of 10 times less than
solar, a star which has [Fe/H] = −2.0 has a metallicity of 100 times less than solar
1

and so on.

1.1

The Milky Way Components and Formation
Scenarios

The Milky Way consists of several different physical components. The populations
of stars that exist in the three main components (the bulge, the disk, and the halo)
have different kinematic and chemical properties. The bulge of the galaxy lies directly
in the center and rotates with a mean velocity of 75 km/s. The bulge has a mean
metallicity of [Fe/H] = +0.25, which is almost twice the solar metallicity. However,
the stars in the bulge span a large range in ages and metallicities (from [Fe/H] = −1.50
to +0.50) and therefore have likely originated from different populations of stars
(Zoccali et al., 2008).
The disk of the Milky Way is composed of several spiral arms and is of course
home to the Sun, which is a distance of approximately 8 kpc away from the center
of our Galaxy. The disk itself can be subdivided into separate components: the thin
disk and the thick disk. The thin disk contains stars which are roughly comparable
in composition to the Sun, having metallicities that are typically within ±2 times the
solar value. The thick disk is slightly more metal-poor overall, having metallicities
that range from −1.6 < [Fe/H] < −0.4.
The halo is the site of some of the oldest and most metal-poor stars in the Milky
Way. The focus of this dissertation is primarily on these halo stars. Like the disk
of the Galaxy, the halo can also be separated into different components which have
different characteristics. The inner halo reaches to 10 − 15 kpc from the Galactic
center. It is slightly oblate, and the stars that populate the inner halo typically have
high orbital eccentricities. The mean metallicity of the inner halo is [Fe/H] = −1.6.
There is also a counter-rotating outer halo component (at distances of above 15 kpc

2

from the Galactic center) which is more spherical and has a mean metallicity of
[Fe/H] = −2.2 (Carollo et al., 2007). The different kinematic properties and chemical
properties of the two halo components are indicative of different structure formation
episodes of the Galactic halo as a whole.
According to current theory, the inner halo was most likely formed by the following
hierarchical scenario. A merger of several sub-Galactic-mass fragments occurred to
form the bulk component of the Milky Way. Subsequent merging of other fragments
resulted in the formation of the inner halo which then necessarily contained the higheccentricity stars that are observed today. As the disk of the Galaxy grew in size,
the surrounding inner halo component was flattened, giving rise to its now-observed
oblate shape (Bekki & Chiba, 2001; Chiba & Beers, 2001).
As the outer halo is both more metal poor and counter rotating, it is likely to have
been formed by another scenario altogether. It is believed that the outer halo was
formed from the accretion of much smaller subsystems, such as the ultra-faint dwarf
galaxies that are observed today. Upon such accretion, these systems would suffer
severe tidal stripping which would result in the field halo stars that are currently
present in the outer halo. Indeed, the recently-studied ultra-faint dwarf galaxies
exhibit the very low metallicities and abundance patterns similar to outer-halo field
stars (Norris et al., 2010a,b).

1.2

The Carbon-Enhanced Metal-Poor Stars of
the Galactic Halo

Metal-poor stars in the Milky Way are archaeological relics of the Galaxy’s formation.
By studying the composition of these objects, it is possible to unearth clues to the
formation and nucleosynthetic history of the Milky Way. Large-scale survey efforts
such as the HK survey (Beers et al., 1985, 1992) and the Hamburg/ESO Survey (HES;

3

Christlieb, 2003; Christlieb et al., 2008) have resulted in the identification of very large
numbers of metal-poor stars (with [Fe/H] < −1.0). A significant fraction of metalpoor stars are found to be enhanced in carbon (Lucatello et al., 2006; Marsteller et al.,
2009), and the fraction of carbon-enhanced metal-poor stars increases with declining
metallicity (Beers & Christlieb, 2005; Frebel et al., 2005; Norris et al., 2007). Much
can be learned from the stellar spectra of carbon-enhanced metal-poor stars.
Carbon-enhanced metal-poor (CEMP) stars preserve chemical signatures of the
Galactic past. By studying their abundance patterns, one can begin to uncover details
of the origins of the elements. There exist a number of different types of CEMP stars
which have unique abundance characteristics. The different classes, as defined by
Beers & Christlieb (2005), are listed in Table 1.1.

Table 1.1: Classification of CEMP Stars

Class

Abundance Pattern

CEMP

[C/Fe] > +1.0

CEMP-r

[C/Fe] > +1.0 and [Eu/Fe] > +1.0

CEMP-s

[C/Fe] > +1.0, [Ba/Fe] > +1.0, and [Ba/Eu] > +0.5

CEMP-r/s

[C/Fe] > +1.0 and 0.0 < [Ba/Fe] < +0.5

CEMP-no

[C/Fe] > +1.0 and [Ba/Fe] < 0.0

The different classes of CEMP stars are suggestive of different sites of carbon production at early times. Often, CEMP stars exhibit enhancements of neutron-capture
elements. The most common types of carbon-enhanced metal-poor stars observed
are the CEMP-s stars, which show evidence of s-process-element enhancements. It is
widely believed that these objects are the result of mass transfer from a companion
low-metallicity asymptotic giant branch (AGB) star, where the production of carbon
4

and s-process-elements occurs. Stars which have r-process enhancements (CEMP-r)
might have evolved from the interstellar medium (ISM) that was polluted by early
Type II supernovae where the r-process is believed to have occurred. There are even
some CEMP stars which have both s- and r-process element enhancements, suggesting that the source of their carbon-enhancements arises from more than one site
and/or time period.
While there are no true “s- or r-process elements”, there are certain isotopes of
elements which are exclusive to each process. Figure 1.1 shows a portion of the path
of the s- and r-processes. Also in Figure 1.1 is the relative fraction of each element
that arises from either process.

Figure 1.1: Part of the paths of the s- and r-processes. The columns on the right
designate the percentage of each element which is populated by the occurrence of
each process. This figure is taken from Sneden et al. (2008). For interpretation of the
references to color in this and all other figures, the reader is referred to the electronic
version of this dissertation.

5

For the CEMP-no class, which exhibit NO neutron-capture element enhancements,
the source of carbon enhancement is less certain. Current theories include the possibility that very massive, rapidly-rotating, mega metal-poor stars (with [Fe/H] < −6.0)
were very efficient producers of carbon, nitrogen, and oxygen (Hirschi et al., 2006;
Meynet et al., 2006). Yet another possible source of CNO enhancement for the observed CEMP-no stars are faint supernovae which undergo heavy mixing and fallback
during their explosions (Umeda & Nomoto, 2003, 2005; Tominaga et al., 2007). The
majority of CEMP-no stars tend to be some of the most metal-poor objects observed.
In fact, they are likely the preservers of the earliest form of carbon production that occurred in the first generations of stars. This idea is confirmed by the recent discovery
of a carbon-enhanced, extremely metal-poor damped Ly-α system at a redshift of z
= 2.34 (Cooke et al., 2011). The abundances determined by this study are consistent
with observed CEMP-no stars in the Galactic halo.
By estimating carbon, nitrogen, oxygen, and neutron-capture-element abundances,
one can classify CEMP stars and begin to constrain the properties of Galactic Chemical Evolution (GCE). Constraints can also be applied to the Initial Mass Function
(IMF) of the Galaxy, as information is revealed about the distribution of different
progenitors of CEMP stars and their masses.
The organization of this dissertation is as follows. Chapter 2 is a version of the
published paper of Kennedy et al. (2011, reproduced by permission of the AAS), and
reports new oxygen abundances for known CEMP stars. Chapters 3 and 4 describe a
new, highly-successful technique to discover previously-unidentified Galactic CEMP
stars. The determinations of C, N, O, and Ba abundances for CEMP stars with the
new X-Shooter spectrograph are presented in chapter 5, and the conclusions related
to all of these studies are discussed in chapter 6.

6

Chapter 2:
Oxygen Abundances for CEMP
Stars from Near-Infrared
Spectroscopy with
SOAR/OSIRIS
2.1

Introduction

Carbon-enhanced metal-poor (CEMP) stars are quite common in the halo populations
of the Milky Way, and are of particular interest, as they preserve important astrophysical information concerning the early chemical evolution of the Galaxy (Beers
& Christlieb, 2005). Previous work has indicated that at least 20% of stars with
metallicities [Fe/H] < −2.0 exhibit large overabundances of carbon ([C/Fe] > +1.0;
Lucatello et al., 2006; Marsteller et al., 2009), although recent studies (e.g. Cohen et
al., 2005; Frebel et al., 2006), have claimed that this fraction is somewhat lower (9%
and 14%, respectively, for [Fe/H] < −2.0). In any case, the fraction of CEMP stars
rises to 30% for [Fe/H] < −3.0, 40% for [Fe/H] < −3.5, and 100% for [Fe/H] < −4.0
(Beers & Christlieb, 2005; Frebel et al., 2005; Norris et al., 2007).
There exist a number of classes of CEMP stars, some of which have been associated with proposed progenitor objects. CEMP-s stars (those with s-process-element
enhancement), for example, are the most commonly observed type to date. Highresolution spectroscopic studies have revealed that around 80% of CEMP stars exhibit s-process-element enhancement (Aoki et al., 2007). The favored mechanism
7

invoked to account for these stars is mass transfer of carbon-enhanced material from
the envelope of an asymptotic giant branch (AGB) star to its binary companion; it
is this surviving binary companion that is now observed as a CEMP-s star.
The class of CEMP-no stars (which exhibit no strong neutron-capture-element
enhancements) is particularly prevalent among the most metal-poor stars. Possible
progenitors for this class include massive, rapidly rotating, mega metal-poor ([Fe/H]
< −6.0) stars, which models suggest have greatly enhanced abundances of CNO due
to distinctive internal burning and mixing episodes, followed by strong mass loss
(Meynet et al., 2006; Hirschi et al., 2006; Meynet et al., 2010). Another suggested
mechanism is pollution of the interstellar medium by the so-called faint supernovae
associated with the first generations of stars, which experience extensive mixing and
fallback during their explosions (Umeda & Nomoto, 2003, 2005; Tominaga et al.,
2007); high [C/Fe] and [O/Fe] ratios are predicted in the ejected material. This model
well reproduces the observed abundance pattern of the CEMP-no star BD+44:493,
the ninth-magnitude [Fe/H] = −3.7 star (with [C/Fe]= +1.3, [N/Fe] = +0.3, and
[O/Fe] = +1.6) discussed by Ito et al. (2009).
The great majority of known CEMP stars were originally identified as metal-poor
candidates from objective-prism surveys, such as the HK Survey (Beers et al., 1985,
1992), and the Hamburg/ESO Survey (HES; Christlieb, 2003; Christlieb et al., 2008),
based on a weak (or absent) Ca II K line. Some candidate CEMP stars also come
from a list of HES stars selected from the prism plates based on their strong molecular
lines of carbon (Christlieb et al., 2001). Medium-resolution spectra for most of these
objects have been obtained over the past few years (Goswami et al., 2006, 2010;
Marsteller, 2007, T. Sivarani et al. 2011, in preparation). Inspection of these data
indicates that at least 50% of these targets are consistent with identification as CEMP
stars, while the others are roughly solar-metallicity carbon-rich stars. Dedicated
surveys for CEMP stars covering a wide range of carbon abundance and metallicities

8

are just now getting underway, based on the observed strength of the CH G band
measured from the HES prism plates (e.g., Placco et al., 2010).
In order to more fully test the association of CEMP-no stars with massive primordial stars and/or faint supernovae, and to better explore the nature of the s-process
in low-metallicity AGB stars (which is still rather poorly understood; Herwig, 2005),
we require measurements of the important elements C, N, and O for as large a sample
of CEMP stars as possible. While estimates of carbon and nitrogen abundances can
be determined from medium-resolution optical or near-ultraviolet spectra of CEMP
stars (e.g., Rossi et al., 2005; Beers et al., 2007b; Johnson et al., 2007; Marsteller
et al., 2009), high-resolution spectroscopy is usually required in order to obtain estimates of oxygen abundances from the forbidden [O1] λ6300 line, the λ7700 triplet
(e.g., Schuler et al., 2006; Sivarani et al., 2006; Fabbian et al., 2009, and references
therein), or the OH lines at 1.5−1.7 µm (Mel´ndez & Barbuy, 2002). Masseron et al.
e
(2010) provide a useful compilation of known elemental abundances for CEMP stars.
In addition to abundance measurements for metal-poor halo stars, oxygen abundances
have also been measured directly in the gas phase in damped Lyα systems (Pettini
et al., 2002, 2008).
If a star has a measured carbon abundance (and, assuming C/O > 1, which applies
for most CEMP stars), essentially all of the O is locked up in CO molecules, and
medium-resolution spectroscopy of the CO ro-vibrational bands in the near-infrared
(near-IR) can be used for the estimation of [O/Fe] (e.g., Beers et al., 2007b, and
references therein). Although one sacrifices measurement accuracy, relative to highresolution studies, this approach has the great advantage that medium-resolution
spectroscopy can be gathered far faster than high-resolution spectroscopy, ensuring
that much larger samples of stars can be investigated. In addition, the large separation
of the 13 CO lines from the 12 CO lines at 2.3µm provides a straightforward means to
measure the important mixing diagnostic 12 C/13 C, as long as the signal-to-noise ratio

9

(S/N) of the spectra is sufficient.
This paper is outlined as follows. In Section 2.2, we discuss details of the observations and data reduction procedures used in the present study. Section 2.3 describes
the previously determined atmospheric parameter estimates and their origins, as well
as details about the synthetic spectra. Methods used for the determination of [O/Fe]
for our sample of stars are described in Section 2.4. Our results and a statistical comparison to high-resolution estimates of [C/Fe] and [O/Fe] for a subset of our program
stars can be found in Section 2.5. Section 2.6 is a short discussion of our results;
conclusions follow in Section 2.7.

2.2

Observations and Data Reduction

Our sample of 57 stars was selected from the HES, based on follow-up mediumresolution optical spectra obtained during the course of searches for low-metallicity
stars. These optical spectra were obtained with the GOLDCAM spectrograph on
the KPNO 2.1m telescope and with the RC Spectrographs on the 4m KPNO and
CTIO telescopes (see Beers et al., 2007b, hereafter Paper I). Additional targets were
selected from the list of carbon-rich candidates published by Christlieb et al. (2001)
with available optical spectra. Based on the optical spectra, all of the candidates are
metal-poor stars, spanning the metallicity range −2.8 ≤ [Fe/H] ≤ −1.0. All of the
stars were selected to be carbon-rich, with the majority exhibiting [C/Fe] ≥ +1.0,
and thus are carbon enhanced as defined by Beers & Christlieb (2005). Since our
intention was to obtain near-IR spectroscopy of the CO features, the stars were also
selected to have effective temperatures less than 5000 K, since warmer stars do not
exhibit strong CO.
Estimates of [O/Fe] for our program stars are derived from the analysis of mediumresolution near-IR spectra taken with the SOAR 4.1m telescope, using the OSIRIS
(Ohio State Infrared Imager/Spectrometer; Depoy et al., 1993) spectrograph during
10

2005 October to 2008 June. We used the long slit (width set to 1”) and long camera
(with focal ratio f/7), which provided a resolving power R = 3000. The long-pass
K-band filter was used to isolate the spectral region from 2.25µm to 2.45µm. Visible
in this band are the four ro-vibrational CO features used for the determination of
[O/Fe]. We also observed A0-type stars at the same air mass as the observations
of the program objects in order to correct for the presence of telluric lines in the
spectra. The K-band magnitude range for our sample stars is ∼ 7 − 12, resulting
in exposure times in the range 600−2400 seconds in order to reach our targeted
S/N of 50/1. Spectra of Ar−Ne arc lamps, taken before or after each program star,
were used for the wavelength calibration of our sample. Bias correction, flat-fielding,
spectral extraction, wavelength calibration, telluric feature correction, and continuum
normalization were all performed using standard IRAF packages. IRAF is distributed
by the National Optical Astronomy Observatory, which is operated by the Association
of Universities for Research in Astronomy, Inc., under cooperative agreement with the
National Science Foundation.

2.3

Adopted Atmospheric Parameters and Synthetic Spectra

Atmospheric parameters (Teff , log g, and [Fe/H]) were estimated from available optical and near-IR photometry as well as from previously obtained medium-resolution
optical spectroscopy. Estimates of Teff are obtained from measured V − K colors
(taken from Beers et al., 2007a, and references therein, as well as from the 2MASS
Point Source Catalog; Skrutskie et al. 2006). The use of near-IR photometry provides for a more accurate determination of Teff , as the K band is less influenced by
the presence of carbon features than bluer bands. We used the Alonso et al. (1996)
calibrations of Teff with V − K colors, as described in Paper I. Surface gravities,

11

log g, have been estimated based on the Padova evolutionary tracks for metallicities
[Fe/H]= −2.5 and [Fe/H]= −1.7 (Girardi et al., 2000; Marigo et al., 2001). Uncertainties in Teff and log g are 100 K and 0.5 dex, respectively. The microturbulence
is taken to be 2 km s−1 for all stars. This is consistent with previously determined
microturbulence values for giant CEMP stars (Johnson et al., 2007; Aoki et al., 2007).
We have constructed two sets of synthetic spectral templates, covering the optical
and near-IR bands. Each set consists of 2000 synthetic spectra with carbon-enhanced
atmospheres generated with the MARCS code (Gustafsson et al., 2008). We used a
previous generation of models here, as updated CEMP models were not available.
We do not, however, anticipate large differences in the models of spectra. The use of
carbon-enhanced models is of particular importance for cool CEMP stars, for which
the atmospheric structure is significantly altered by carbon (Masseron et al., 2006).
No 3D→1D corrections have been applied to our estimates. Recent studies of these
effects on two hyper metal-poor stars (Collet et al., 2006) have revealed [O/Fe] corrections of ∼ −0.8 based on OH molecules, thereby lowering the measured abundance
of oxygen. However, the magnitude of such corrections is expected to decrease with
increasing metallicity (Collet et al., 2007). As the metallicities of our targets range
from −1.0 to −2.8, it is likely that the three-dimensional effects are less severe. In
addition, more recent studies have been carried out (A. Ivanauskas, private communication) concerning the effects of convection on C2 , CH, CN, CO, NH, and OH
molecules. When compared to the results from Collet et al. (2006), the magnitude of
the corrections appears smaller.
The synthetic grid covers a range Teff = 4000 to 6000 K , log g = 0.0 to 5.0, [Fe/H]
= −5.0 to 0.0, and [C/Fe] = 0.0 to +4.0. We adopt fixed nitrogen abundances set
to 0.5 dex less than the carbon abundances, which is roughly appropriate for CEMP
stars. The CH and CN line lists used for the synthesis of the optical spectra are those
compiled by Plez (see Plez & Cohen, 2005). The CO line lists used for the near-IR

12

synthesis are taken from Kurucz (1993). The synthetic grids are then degraded to
match the resolving power of the observed spectra (R = 2000 for the optical spectra
and R = 3000 for the near-IR spectra).
The optical spectra are used for the determination of [Fe/H] and [C/H]. The Ca
II K line is matched with the model spectra to estimate [Fe/H], and the C2 and CN
features are fit for the estimation of [C/H] (see Paper I). Our adopted atmospheric
parameters, as well as the derived [C/H] and [C/Fe], are listed in Tables 2.1, 2.2, and
2.3 .
Table 2.1: Atmospheric Parameters and Carbon Abundances I

Name

T eff (K)

log g (cgs)

[Fe/H]

[C/H]

[C/Fe]

HE 0002+0053

4225

0.27

−2.18

−0.03

2.15

HE 0010−3051

4177

0.17

−2.71

−0.40

2.31

HE 0017+0055

4185

0.18

−2.72

−0.41

2.31

HE 0033−5605

4021

0.00

−1.48

−0.26

1.22

HE 0043−2433

4397

0.61

−1.47

−0.32

1.15

HE 0111−1346

4651

1.08

−1.91

−0.22

1.70

HE 0120−5834

4828

1.62

−2.40

0.00

2.40

HE 0140−3956

4468

0.84

−2.04

−0.23

1.81

HE 0151−0341

4849

1.42

−2.46

0.00

2.46

HE 0155−2221

4109

0.00

−2.28

−0.50

1.78

HE 0206−1916

4741

1.23

−2.83

−0.50

2.33

HE 0219−1739

4227

0.27

−1.50

0.00

1.50

HE 0251−2118

4710

1.16

−1.50

−0.37

1.13

HE 0310+0059

4861

1.69

−1.32

−0.07

1.24

HE 0314−0143

4201

0.22

−1.25

−0.36

0.89

HE 0319−0215

4416

0.64

−2.42

−0.33

2.09

HE 0330−2815

4411

0.64

−1.46

0.00

1.46

HE 0359−0141

4340

0.54

−1.73

−1.00

0.73

13

Table 2.2: Atmospheric Parameters and Carbon Abundances II

Name

T eff (K)

log g (cgs)

[Fe/H]

[C/H]

[C/Fe]

HE 0408−1733

4260

0.33

−2.06

−1.00

1.06

HE 0417−0513

4669

1.22

−1.88

−1.00

0.88

HE 0419+0124

4368

0.61

−1.49

−1.00

0.49

HE 0429+0232

4409

0.63

−2.05

−1.00

1.05

HE 0430−1609

4651

1.08

−2.14

−0.33

1.81

HE 0439−1139

4833

1.62

−1.31

−0.50

0.81

HE 0457−1805

4484

0.77

−1.46

−0.46

0.99

HE 0458−1754

4374

0.56

−1.95

−1.00

0.95

HE 0507−1653

4880

1.50

−1.81

−0.04

1.77

HE 0518−1751

4252

0.32

−1.90

−1.00

0.90

HE 0519−2053

4775

1.46

−1.45

−0.50

0.95

HE 0547−4428

4217

0.25

−1.97

−0.40

1.57

HE 1011−0942

3716

5.00

−1.29

−0.41

0.88

HE 1023−1504

4421

0.66

−2.50

−0.06

2.44

HE 1125−2942

3947

5.00

−1.01

0.00

1.01

HE 1145−0002

4033

0.00

−1.49

−0.18

1.31

HE 1152−0355

4214

0.23

−1.74

−0.15

1.59

HE 1152−0430

4573

0.93

−1.50

−0.17

1.33

HE 1204−0600

4581

1.07

−1.50

−0.02

1.48

HE 1207−3156

4664

1.23

−2.03

−0.50

1.53

HE 1230−0230

3966

5.00

−1.15

−0.30

0.85

HE 1238−0435

4433

0.68

−2.27

−0.39

1.88

HE 1246−1510

3776

5.00

−1.50

−0.35

1.15

HE 1255−2324

4405

0.69

−1.47

0.00

1.47

HE 1331−0247

4238

0.28

−1.76

−0.25

1.52

14

Table 2.3: Atmospheric Parameters and Carbon Abundances III

Name

log g (cgs)

[Fe/H]

[C/H]

[C/Fe]

HE 1410−0125

4427

0.66

−2.20

−0.50

1.70

HE 1418+0150

4163

0.13

−2.34

−0.50

1.84

HE 1428−1950

4531

0.85

−2.07

−0.28

1.78

HE 1431−0755

4104

0.00

−2.11

−0.01

2.10

HE 1524−0210

4145

5.00

−1.50

−0.50

1.00

HE 2115−0522

4762

1.46

−1.50

−1.00

0.50

HE 2145−1715

4461

0.77

−1.50

−0.36

1.14

HE 2153−2323

4261

0.33

−2.16

−0.50

1.66

HE 2200−1652

4473

0.84

−1.50

0.00

1.50

HE 2207−1746

4737

1.23

−1.87

−0.27

1.60

HE 2221−0453

4129

0.04

−2.58

−0.50

2.08

HE 2224−0330

4831

1.62

−1.20

−0.47

0.73

HE 2228−0137

4368

0.61

−2.39

−0.22

2.17

HE 2339−0837

2.4

T eff (K)

4939

1.60

−2.71

0.00

2.71

Determination of [O/H]

In order to determine [O/H], we employed the near-IR synthetic spectra constructed
from model atmospheres with carbon enhancements (see above). Each synthetic
spectrum covers the wavelength range 2.25µm−2.40µm. For all combinations of these
parameters, models were available with [O/Fe] values of 0.0, +0.4, and +0.8.
The first step in the determination of [O/Fe] is to use the grid of synthetic spectra
in combination with the atmospheric parameters to create three models with [O/Fe]
15

values of 0.0, +0.4, and +0.8. These models are then used in order to estimate the
[O/Fe] of the program stars. Each star in the sample has a previously estimated
Teff , log g, [Fe/H], and [C/Fe]. Given the fact that four parameters are known, the
routine begins with 48 models: 16 models for each value of [O/Fe]. The 48 models
are selected as having the two closest values of Teff , log g, [Fe/H], and [C/Fe] for
each of the three values of [O/Fe]. Once selected, a linear interpolation over each
parameter is performed in order to create three final models, one for each value of
[O/Fe], with the known values of the four parameters. The three final models for
a set of typical atmospheric parameters are shown in Figure 2.1. With the other
parameters fixed, it is easy to see how a typical spectrum changes with increasing
oxygen abundance. Next, a linear interpolation over [O/Fe] is performed on the three
final models, creating new model spectra with varying oxygen. It should be noted
that while the interpolation was performed over [O/Fe], the metallicity is fixed, and
therefore the oxygen varies as [O/H]. In some cases, it was necessary to extrapolate
beyond the boundaries of the model grid in order to find a good fit to the data.
Due to the presence of four large CO bands in the near-IR spectra, there were
difficulties in fitting the continuum across the entire region. It was important to fit
the continua with a low-order function so that the depth of the absorption features
was not artificially enhanced or lessened due to the continuum fit. For this reason,
the spectra of all stars in the sample were trimmed around each of four CO bands,
and a local continuum was fit for each band prior to spectral synthesis. With the
use of the synthetic spectra, oxygen abundances were then estimated individually for
each of the four bands by minimizing χ2 . A robust average using bisquare weighting
of the four separate values was taken as the final estimate of oxygen abundance, with
an associated robust estimate of the scatter in these values taken as the error of
determination.
Figure 2.2 shows the fitting technique applied to four stars from the sample. Each

16

Figure 2.1: Three synthetic spectra with different [O/Fe] ratios. Each spectrum has
Teff of 4500 K, log g of 1.0, [Fe/H] of −2.0, and [C/Fe] of +1.0.
row shows the four separate estimates of [O/H] for each star. Also plotted on each
panel are synthetic spectra with [O/H] values that vary from the best-fitting spectra
by ±0.5 dex. Once the robust average is applied the resulting estimates of [O/H] are
−1.2 for HE 0111−1346, −0.5 for HE 0519−2053, −0.6 for HE 1207−3156, and −1.2
for HE 2339−0837.

2.5

Results

The distributions of [O/Fe] versus three of the parameters used for their determinations are shown in Figure 2.3. The solid lines are linear fits to the data. For the entire
sample, there are no significant correlations of [O/Fe] with Teff , [Fe/H], or [C/Fe]. In
the middle panel of Figure 2.3, it can be seen that only one of the stars in our sample
with [Fe/H] < −2.5 has a value of [O/Fe] less than +1.0. The bottom panel of Figure

17

Figure 2.2: Each row shows four estimates of [O/H] for a star from our sample:
HE 0111−1346, HE 0519−2053, HE 1207−3156, and HE 2339−0837, respectively. In
each panel, the black lines are the data, the red lines are the best-fitting synthetic
spectra, the green lines have [O/H] values of 0.5 dex lower than the best-fitting
spectra, and the blue lines have [O/H] values of 0.5 dex higher than the best-fitting
spectra. A robust average of the four separate estimates is taken as the final estimate
of oxygen abundance for each star. See the text for details.
2.3 shows the distribution of [O/Fe] versus [C/Fe], revealing that the majority of our
sample (45 stars) have [C/Fe] > +1.0, and so meet the definition for CEMP stars
given by Beers & Christlieb (2005). The other stars exhibit carbon enhancements of
[C/Fe] ≥ +0.5 and thus are at least moderately enhanced in carbon. The average
error in the determination of [O/Fe] for our entire sample is 0.4 dex. We adopted a
minimum error for our [O/Fe] estimates of 0.25 dex due to the influence of errors that
arise from the estimation of Teff , log g, [Fe/H], and [C/Fe] (see Paper I for details).
In Figure 2.4, a carbon cut has been made, such that only the oxygen abundances

18

Figure 2.3: Top panel: [O/Fe] vs. Teff for the entire sample. Middle panel: [O/Fe] vs.
[Fe/H] for the entire sample. Bottom panel: [O/Fe] vs. [C/Fe] for the entire sample.
for those stars with [C/Fe] > +1.75 are plotted against [Fe/H]. The stars with the
highest abundances of carbon exhibit some of the lowest metallicities in our sample.
This is not surprising, given that high values of [C/Fe] are often associated with
lower metallicities. The fit to the data shows a slight trend of increasing oxygen with
decreasing metallicity. The solid line is a least squares fit of [O/Fe] as a function
of [Fe/H]. Only a marginally significant slope (−0.616 ± 0.314) is found, hence the
correlation is quite weak. For comparison, the dashed line in this figure represents
the fit of [O/Fe] versus [Fe/H] for the carbon-normal stars from the Spite et al.
(2005) sample. The [O/Fe] estimates of the Spite et al. (2005) sample come from the
forbidden O I λ6300 line.
19

Figure 2.4: [O/Fe] vs. [Fe/H] for the stars with [C/Fe] ≥ +1.75. The dashed line
represents the fit for the carbon-normal stars from the Spite et al. (2005) sample of
very metal-poor stars, while the solid line is the best fit for our data.

2.5.1

Statistical Comparison to High-resolution Oxygen
Estimates

The present catalog of measured oxygen abundances available in the literature for
CEMP stars is still relatively small, due to the difficulty of obtaining estimates from
optical spectra, even at high spectral resolution. However, we can at least compare
the regions of the [O/Fe] versus [Fe/H] parameter space that are occupied by CEMP
stars of various sub-classes, based on previous high-resolution oxygen estimates, with
those from our present medium-resolution effort.
Figure 2.5 shows [O/Fe] for our entire sample, with different boxes indicating the
regions of parameter space occupied by several classes of CEMP stars. Sources for

20

the high-resolution data for different classes of CEMP stars can be found in Masseron
et al. (2010), and references therein. The majority of the stars in our sample occupy
regions of the diagram as CEMP stars that have confirmed, high-resolution measurements of s-process-element enhancements. There is clearly overlap with the region
occupied by CEMP-r/s stars as well. Few of our stars overlap with the region occupied
by CEMP-no stars in the literature; CEMP-no stars tend to be more metal-deficient
than most of the stars in our sample.

Figure 2.5: [O/Fe] vs. [Fe/H] for the entire sample. The colored boxes show the
regions occupied by different types of CEMP stars found in Masseron et al. (2010).

2.5.2

High-resolution Nitrogen Estimates

For 13 of our program stars, high-resolution estimates of [N/Fe] are available from S.
Lucatello (private communication) and/or Aoki et al. (2007). We selected 10 of these
21

stars for which the available high-resolution estimates of [C/Fe] were within 0.5 dex
of our medium-resolution estimates. The three that are omitted from our analysis
and discussion have associated high-resolution [Fe/H] and/or [C/H] estimates that
differ significantly from the medium-resolution estimates of these species. We report
values of high-resolution [C/Fe], [C/Fe]h , using our estimates of [Fe/H] combined with
the high-resolution [C/H]. We report high-resolution estimates of [N/Fe], [N/Fe]h , by
combining our estimates of [Fe/H] with the high-resolution values of [N/H]. Two of
the 10 stars had high-resolution estimates both from Lucatello and Aoki et al. (2007),
and an average of the two available estimates was taken.
The values of [C/Fe], [C/Fe]h , [N/Fe]h , [O/Fe], and σ [O/Fe] for our entire sample
are listed in Tables 2.4, 2.5, and 2.6.

Table 2.4: Abundance Ratios for the Entire Sample I

Name

[C/Fe]

[C/Fe]h

[N/Fe]h

[O/H]

[O/Fe]

σ([O/F e])

HE 0002+0053

2.15

...

...

−1.6

0.6

0.3

HE 0010−3051

2.31

...

...

−1.7

1.0

0.7

HE 0017+0055

2.31

1.82

0.52

−1.7

1.0

0.4

HE 0033−5605

1.22

...

...

−0.5

1.0

0.7

HE 0043−2433

1.15

...

...

−0.6

0.8

0.5

HE 0111−1346

1.70

1.48

0.88

−1.2

0.8

0.3

HE 0120−5834

2.40

...

...

−1.7

0.7

0.3

HE 0140−3956

1.81

1.55

1.15

−1.4

0.7

0.3

HE 0151−0341

2.46

2.16

2.06

−1.3

1.1

0.3

HE 0155−2221

1.78

...

...

−1.4

0.9

0.8

HE 0206−1916

2.33

2.42

1.92

−1.7

1.1

0.3

HE 0219−1739

1.50

...

...

−1.9

−0.4

0.4

HE 0251−2118

1.13

...

...

−0.9

0.6

0.4

HE 0310+0059

1.24

...

...

−0.6

0.7

0.6

22

Table 2.5: Abundance Ratios for the Entire Sample II

Name

[C/Fe]

[C/Fe]h

[N/Fe]h

[O/H]

[O/Fe]

σ([O/F e])

HE 0314−0143

0.89

...

...

−0.3

0.9

0.9

HE 0319−0215

2.09

2.12

0.92

−2.1

0.4

0.5

HE 0330−2815

1.46

...

...

−0.7

0.8

0.6

HE 0359−0141

0.73

...

...

−0.8

0.9

0.6

HE 0408−1733

1.06

...

...

−1.3

0.8

0.4

HE 0417−0513

0.88

...

...

−0.3

1.6

0.5

HE 0419+0124

0.49

...

...

−0.9

0.6

0.4

HE 0429+0232

1.05

...

...

−1.6

0.4

0.3

HE 0430−1609

1.81

1.84

0.64

−2.1

0.0

0.3

HE 0439−1139

0.81

...

...

0.4

1.8

0.3

HE 0457−1805

0.99

...

...

−1.1

0.4

0.4

HE 0458−1754

0.95

...

...

−1.8

0.2

0.3

HE 0507−1653

1.77

1.61

0.97

−1.4

0.4

0.3

HE 0518−1751

0.90

...

...

−1.9

0.1

0.3

HE 0519−2053

0.95

...

...

−0.5

1.0

0.3

HE 0547−4428

1.57

...

...

−1.1

0.9

0.7

HE 1011−0942

0.88

...

...

−0.3

1.0

0.3

HE 1023−1504

2.44

...

...

−1.2

1.3

0.3

HE 1125−2942

1.01

...

...

0.1

1.2

0.3

HE 1145−0002

1.31

...

...

−1.3

0.2

0.4

HE 1152−0355

1.59

...

...

−1.4

0.4

0.6

HE 1152−0430

1.33

...

...

−0.3

1.2

1.1

HE 1204−0600

1.48

...

...

−0.4

1.1

0.4

HE 1207−3156

1.53

...

...

−0.6

1.5

0.3

HE 1230−0230

0.85

...

...

0.5

1.7

0.4

23

Table 2.6: Abundance Ratios for the Entire Sample III

Name

[C/Fe]h

[N/Fe]h

[O/H]

[O/Fe]

σ([O/F e])

HE 1238−0435

1.88

...

...

−1.6

0.6

0.5

HE 1246−1510

1.15

...

...

0.1

1.6

0.3

HE 1255−2324

1.47

...

...

−1.6

−0.1

0.3

HE 1331−0247

1.52

...

...

−1.3

0.5

0.4

HE 1410−0125

1.70

...

...

−1.4

0.8

0.4

HE 1418+0150

1.84

...

...

−1.1

1.3

0.5

HE 1428−1950

1.78

...

...

−1.1

1.0

0.3

HE 1431−0755

2.10

...

...

−0.8

1.3

0.3

HE 1524−0210

1.00

...

...

0.4

1.9

0.3

HE 2115−0522

0.50

...

...

−0.5

1.0

0.8

HE 2145−1715

1.14

...

...

−0.7

0.8

0.7

HE 2153−2323

1.66

...

...

−1.7

0.5

0.4

HE 2200−1652

1.50

...

...

−1.7

−0.2

0.3

HE 2207−1746

1.60

...

...

−0.9

1.0

0.3

HE 2221−0453

2.08

2.08

1.19

−2.1

0.5

0.3

HE 2224−0330

0.73

...

...

−0.8

0.4

0.6

HE 2228−0137

2.17

1.90

1.60

−1.9

0.5

0.4

HE 2339−0837

2.6

[C/Fe]

2.71

...

...

−1.2

1.5

0.3

Discussion

We expect that the majority of CEMP stars in our sample have been polluted by a
companion low-metallicity AGB star. In an AGB star, intershell oxygen is predicted
24

to be closely related to intershell 12 C, which in turn has a direct influence on the
maximum 13 C abundance. By studying these abundances, we can better understand
the nature of the s-process, as the 13 C(α, n)16 O reaction is a major neutron source
for the s-process in AGB stars (Lugaro et al., 2003). According to current theory,
oxygen production in AGB stars becomes increasingly significant with decreasing
metallicity (Herwig, 2004, 2005; Campbell & Lattanzio, 2008; Lau et al., 2009). Since
the overabundance of oxygen is smaller at solar metallicities, lower metallicities are
better suited for probing the primary production of oxygen. In addition, the observed
abundance patterns of elements produced by the progenitor are expected to depend
on the mass (and metallicity) of the AGB star (Herwig, 2004; Stancliffe & Glebbeek,
2008).
In the following subsections, we consider a number of issues that could potentially
impact the interpretation of our measurements. First, we compare the results of our
CEMP stars to those of carbon-normal metal-poor stars (Section 2.6.1). We then turn
to the recent literature on the abundance yields of low-metallicity AGB models to
compare with our derived C, N, and O abundances (Section 2.6.2). We also consider
how dilution processes might be expected to alter our abundances (Section 2.6.3).
Finally, we address the sources of uncertainty in the AGB models, and how these
may lead to altered abundances of C, N, and O (Section 2.6.4).

2.6.1

[O/Fe] in Carbon-normal and Carbon-enhanced Metalpoor Stars

The linear fit to the full sample shown in the middle panel of Figure 2.3 is consistent
with the fit from the Spite et al. (2005) sample. However, most of the values of
[O/Fe] from the carbon-normal sample are tightly distributed around a constant value
of +0.7. Referring back to Figure 2.5, we see much more scatter in our values of
[O/Fe] for the carbon-enhanced stars, with values reaching as high as +2.0. It can be
25

inferred that metal-poor stars, regardless of carbon enhancement, commonly exhibit
enhancements of oxygen. However, when carbon enhancement is present, additional
oxygen enhancement can be expected as well, due to the fact that both of these
elements are enhanced by some of the same mechanisms.

2.6.2

C, N, and O: Comparison with AGB Models

For 10 of the 13 stars for which we report high-resolution [N/Fe] estimates in Tables
2.4, 2.5, and 2.6, Figure 2.6 shows [C/Fe], [N/Fe], and [O/Fe] as a function of [Fe/H].
One can notice an increase in the abundances of all three species with decreasing
metallicity. The linear fits for each of the species all have approximately the same
slope, suggesting that abundances of carbon, nitrogen, and oxygen are highly correlated with one another. With estimates of [C/Fe], [N/Fe], and [O/Fe], we can compare
our results to the predictions of abundance yields due to AGB evolution as described
by Herwig (2004). Both carbon and oxygen are dredged up to the surface in AGB
stars after thermal pulses. The overabundance of these elements in low-metallicity
AGB stars is larger for lower initial masses (Herwig, 2004), due to the fact that the
intershell mass is larger. With a similarly large dredge up parameter and a smaller
envelope, the enrichment of C and O in the envelope of a lower mass star is larger.
This can be seen in the top panel of Figure 2.7, where the abundance predictions
for C, N, and O (Herwig, 2004) are shown for five different AGB masses, ranging
from 2 M⊙ to 6 M⊙ , all with [Fe/H] = −2.3. The C, N, and O abundances for 10
stars with available [N/Fe] in our sample are shown in a similar way in the lower two
panels of Figure 2.7. For most of the stars, the relationship of carbon and oxygen is
consistent with the models. The abundance pattern for HE 0017+0055 is a very close
match to an AGB star of about 3 M⊙ . The other nine stars exhibit some discrepancy
with respect to nitrogen, which has been noted before for other CEMP stars with
s-process-element enhancement (Paper I; Sivarani et al., 2006). A previous effort to

26

search for metal-poor stars with large enhancements of nitrogen relative to carbon
yielded similar abundances (Johnson et al., 2007). All but five of the stars in our sample are giants, and thus more mixing and dilution of any material transferred from
an AGB companion is expected, thereby resulting in such intermediate abundances
of nitrogen (Paper I; Denissenkov & Pinsonneault, 2008). In addition, the possible
occurrence of H-ingestion flashes (HIFs; Herwig, 2003, 2005; Woodward et al., 2008;
Hajduk et al., 2005; Campbell & Lattanzio, 2008) could potentially enhance N in
metal-poor stars.

Figure 2.6: C, N, and O abundances vs. metallicity for 10 stars from our sample.
Also shown are linear fits for these species.
Nitrogen is an element that is very sensitive to CN cycling, and the C/N ratio
indicates if the CN cycle has been activated partially, or whether mixing and thermodynamic conditions have been such that the CN cycle has reached equilibrium. The
27

Figure 2.7: Top panel: predicted abundances of C, N, and O for model AGB stars of
different masses (in M⊙ ) from Herwig (2004). Middle and bottom panels: abundances
of C, N, and O for 10 stars from our sample.
latter is the case in hot-bottom burning (HBB), which is found in more massive AGB
stars (Boothroyd et al., 1993). In this case, the bottom of the convective envelope
connects with the H-burning shell, allowing the processing of envelope CN material in
the H shell. Material lost at the surface is accordingly modified in its CN abundance
ratios. The limiting mass for the onset of HBB decreases with decreasing metallicity
from ∼ 5 M⊙ at Z = 0.02 to ≤ 3 M⊙ for Z = 0.0 (Forestini & Charbonnel, 1997;
Siess et al., 2002). Stars which experience HBB show a C/N ratio close to the equi-

28

librium ratio (< 0.1), as can be seen also in Figure 2.7 for the 4, 5, and 6 M⊙ cases.
In contrast, the 2 and 3 M⊙ cases show very large C/N ratios. Not even partial CN
cycling has occurred in these models, and the small overabundance of N is entirely
due to the first and second DUP (Herwig, 2004). Partial CN cycling at the bottom
of low-mass giant envelope convection is a well observed and parametrically modeled
feature (e.g. Denissenkov & VandenBerg, 2003). Whether the same process operates
in AGB stars as well is currently debated (e.g. Karakas et al., 2010).

2.6.3

Considering the Effects of Dilution

Nitrogen estimates only exist for 10 of our stars. We compared the carbon and
oxygen predictions from Herwig (2004) to all of our CEMP stars. None of the stars
with available nitrogen estimates show the extremely low C/N ratio indicative of
HBB. We therefore restrict the comparison of our sample to the 2 and 3 M⊙ models
from Herwig (2004). In Figure 2.8, we show all of our stars that have [C/Fe] ≥ +1.0.
The black square and the red triangle at the upper right of the figure are the model
predictions for 2 M⊙ and 3 M⊙ , respectively. Clearly these predictions have higher
carbon and oxygen estimates than our sample, but this is likely due to the fact that
the effects of dilution are not considered. We consider a parametric mixing model
to test the effects of dilution of the accreted material from an AGB companion. We
chose a range of initial masses of the observed star (0.5 − 0.9 M⊙ ) and accreted
masses (0.1 − 0.5 M⊙ ) and assumed complete mixing. We set the maximum mass of
the observed star at 1 M⊙ . In order to consider only the AGB-phase contributions to
carbon and oxygen, we subtracted off the likely contribution to carbon and oxygen
arising from pre-star formation enhancements. These contributions were chosen to be
the average enhancements of carbon and oxygen from the Spite et al. (2005) sample
of unmixed metal-poor stars. The subtracted enhancement of [C/Fe] was 0.18 ± 0.16
dex and the subtracted enhancement of [O/Fe] was 0.7 ± 0.17 dex. The resulting

29

predicted abundances based on this simple dilution experiment are shown as the
black and the red lines in Figure 2.8. Once the effects of dilution are considered,
the AGB predictions of Herwig (2004) fall within the parameter space of our sample.
The magenta symbols are for the range of metallicity that is most consistent with the
AGB models. For lower metallicities, the AGB model values for [C/Fe] and [O/Fe]
would be larger, and for higher metallicity, they would decrease (Herwig, 2004).

Figure 2.8: [O/Fe] vs. [C/Fe] for all stars from our sample with [C/Fe]≥ +1.0, colorcoded by their metallicity ranges, as noted in the legend. The black square and
red triangle mark the values of the abundances predicted by Herwig (2004) and the
colored lines represent the effects of dilution once the AGB material is accreted onto
the observed CEMP star. See text for more details.

30

2.6.4

Uncertainties of the AGB Models

The models of Herwig (2004) should be considered as rather conservative, standard
predictions that suffer from the usual uncertainties associated mostly with convective
mixing and mass loss. We can obtain an indication of the order of magnitude of these
uncertainties by comparing the Herwig (2004) models with those of other authors.
Comparing the 2 M⊙ case with that of Cristallo et al. (2009a), the carbon predictions
agree very well, while the [N/Fe] prediction of Cristallo et al. (2009a) is about 0.2 dex
higher and their O-overabundance prediction is 0.8 dex lower than that of Herwig
(2004). Karakas (2010) provides a comparison between her Z = 0.0001, 2 M⊙ yield
predictions and that of Cristallo et al. (2009a) which shows that her C, N, and O
yields are all approximately twice those of Cristallo et al. (2009a), which is easiest
to be understood in terms of a lower mass loss rate in the Karakas (2010) yields.
Karakas (2010) provides a more in-depth discussion of model prediction differences
from different authors.
In addition, we have to discuss uncertainties deriving from entirely alternative
evolution scenarios. Here, we should mention the possible occurrence of HIFs, which
have been introduced in Section 2.6.2. These events are also referred to as protoningestion episodes or double He-shell flashes. There are several uncertainties related
to these combustion-type flashes in which protons are convectively mixed into the Heshell flash convection zone, and release energy on the dynamic timescale of convective
flows.
First, the occurrence of these events, in which the He-shell flash convection zone
has to break through the entropy barrier associated with the H-shell, is more likely
with lower CNO abundances in the envelope. For example, an alpha-enhancement of
the initial abundance composition without change of [Fe/H] will make the HIF less
likely or even suppress it (Cristallo et al., 2009a). If the stars are rotating (Meynet
& Maeder, 2002), the CNO abundance may be enhanced during the core He-burning
31

phase, which may also suppress the HIF (Herwig, 2003).
The second uncertainty in HIF predictions is the quantitative mixing and nucleosynthesis in a convective combustion regime that breaks some of the assumptions
of mixing-length theory and one-dimensional spherically symmetric stellar evolution
(Herwig et al., 2010). Keeping these significant uncertainties in mind, we nevertheless
note that HIF models would predict the N signature of partial burning. For example, considering Figure 6 in Cristallo et al. (2009b), the prediction of an HIF model
followed by maybe only a few pulses would correspond well to the CNO abundance
patterns observed in most of our stars shown in Figure 2.7.

2.7

Conclusions

We have used near-IR medium-resolution spectroscopy in order to estimate [O/Fe]
for a sample of candidate carbon-enhanced stars selected from the HES. This method
of abundance analysis allows us to obtain oxygen abundances accurate to about 0.4
dex. The use of four separate CO features to estimate oxygen abundances from the
near-IR spectra allows for more precise estimates, based on a robust average of the
independently determined fits. A large spread of derived [O/Fe] values are obtained
for this sample, ranging from near the solar value to as much as one hundred times
greater.
A comparison of our abundance determinations with high-resolution estimates
was carried out. The values of [O/Fe] for our full set of 57 CEMP stars largely
fall within regions of parameter space occupied by the high-resolution estimates of
oxygen for other CEMP stars. We also found that the majority of our stars have
oxygen abundances that are consistent with known CEMP-s and CEMP-r/s stars.
Only a few stars could be considered CEMP-no stars, based on the data compiled in
Masseron et al. (2010). This is likely due to the fact that CEMP-no stars commonly
have lower metallicities than most of the stars in this sample.
32

Oxygen enhancements (on the order of [O/Fe] = +0.7) have also been observed
in very metal-poor stars without significant carbon enhancement, indicating that
there were early oxygen-producing nucleosynthetic sites in the Galaxy independent
of any enhancement by AGB evolution. However, we find that the [C/Fe], [O/Fe],
and [N/Fe] (when available) estimates follow the patterns from Herwig (2004) closely
enough that mass transfer from an AGB companion is a likely scenario for many of
the stars in our sample, especially when the effects of dilution are considered.
Our measured carbon abundances always exceed the available high-resolution
[N/Fe] abundances. If the origin of CNO abundance patterns comes from HBB in
an intermediate mass (AGB) star, one would expect to see elevated [N/Fe] relative
to [C/Fe] and [O/Fe]. This signature is not found in our sample, but it has been
suggested that other mechanisms, such as cool-bottom-processing (Wasserburg et al.,
1995; Denissenkov & VandenBerg, 2003) or the occurrence of HIF, can alter the levels
of nitrogen enhancement.
It is likely that the majority of CEMP stars in this sample will turn out to be
enhanced in neutron-capture elements. Consistency of most of our program stars
with the CEMP-s class, based both on comparison to AGB models and existing
high-resolution data, is expected since that CEMP stars with s-process-element enhancement are the most commonly observed type to date. However, recent chemical
evolution models (Cescutti & Chiappini, 2010) have revealed that the winds from massive, rapidly rotating metal-poor stars can result in a large scatter in the predicted
abundances of C, N, and O, presumably without the production of neutron-capture
elements. Therefore, we are currently unable to assign classification to this sample
of CEMP stars. High-resolution spectra of the stars in our sample will help clarify
these questions.

33

Chapter 3:
A Survey to Identify New CEMP
Stars
3.1

Motivation

A large fraction of metal-poor stars exhibit enhancements of carbon, and therefore
most of the carbon-enhanced metal-poor (CEMP) stars that have been observed and
studied to date were originally identified with objective-prism surveys like the HK
Survey (Beers et al., 1985, 1992), and the Hamburg/ESO Survey (HES; Christlieb,
2003; Christlieb et al., 2008), based on a weak (or absent) Ca II K line. Efforts
to try to identify new CEMP stars have been successfully implemented in the past
(Christlieb et al., 2001). A subset of candidate carbon-enhanced stars from the HES
was selected based on the sum of the strengths of molecular features of carbon (CN,
CH, and C2 ) in the spectra. Follow-up medium-resolution spectroscopy of these objects was carried out (Goswami et al., 2006; Marsteller, 2007), and it was discovered
that about 50% of the candidates were CEMP stars as defined by Beers & Christlieb
(2005). This selection technique, however, has limitations. The candidates are typically biased towards lower temperatures, as the strength of the molecular carbon
features diminishes with higher temperatures. Often, carbon-enhanced stars with
Teff > 5500 K only exhibit a strong CH G band. In addition, since most CEMP stars
were discovered during searches for metal-poor stars, there are sure to be unidentified
CEMP stars in the intermediate metallicity range (−2.0 < [Fe/H] < −1.0).

34

A new technique has been developed to identify CEMP stars which is based solely
on the strength of the CH G band. This new selection technique eliminates both
the temperature and the metallicity bias that had previously existed for new CEMP
discoveries. It is described in the following sections.

3.2

The Selection Technique: A New Extended
Index for the CH G Band

In the past, carbon abundances were determined by employing a 15 ˚ wide line
A
A
index (GP index) centered at the CH G band at 4304 ˚. These indices are based
on a continuum that is linearly fit between two sidebands on the red and blue side
of the CH feature. Several studies, including Rossi et al. (2005), have shown that
this line index is not wide enough to be fully representative of the strength of the G
band. A new line index, GPHES, with a width of 26 ˚ and wider-set sidebands was
A
developed by Christlieb et al. (2008), but even this was far too narrow to account for
the depression of the continuum around the CH G band.
A new 200 ˚ wide line index centered at 4300 ˚ has been defined in order to
A
A
account for the large absorption features. The continuum is fit over the entire spectral
region, rather than relying on two sidebands for a linear fit. The GPE (GPHES
Extended) line index is defined as:
4400

1−

GP E =
4200

S(λ)
C(λ)

dλ,

where S(λ) is the observed spectrum and C(λ) is the local continuum. Figure 3.1
shows the comparison of the old GPHES index with the new GPE index for a typical
carbon-enhanced metal-poor star. It is clear from this figure that even the GPHES
is too narrow, given that the sidebands can easily fall in regions that are depressed
with respect to the real continuum.
35

1000
HE 1046−1352
GPE = 52.1
GPHES = 8.6
KPHES = 1.9
(J−K)0 = 0.465

800

counts

600

400

200

0
4500

4400

4300

4200

4100
λ(Å)

4000

3900

3800

3700

3600

Figure 3.1: The new GPE index. The solid black line shows the continuum fitting
applied to the spectra. The long-dashed green lines mark the position of the the Ca
II H and K features. The GPHES is denoted by the blue dotted lines, and the newly
defined GPE is denoted with the black dashed lines. The arrows mark the values of
the GPHES continuum sidebands. This figure is from Placco et al. (2010).

3.3

Target Selection

All targets for the study were taken from the Hamburg/ESO Survey (HES). First, only
stars were selected which had accurate 2MASS JHK photometry, and we restricted
the color range to be 0.15 ≤ (J−K)0 ≤ 0.90. Next, we selected targets which had
a KPHES value of less than 8.0. The KPHES (Christlieb et al., 2008) is a similar
line index to the KP index defined by Beers et al. (1985, 1999), whereby metallicity
can be determined by the strength of the Ca II K line. By selecting on targets with
KPHES less than a value of 8.0, we eliminated the stars with Ca II K lines that are
too strong to be indicative of a metal-poor star (with [Fe/H] < -1.0). We also selected
only stars with B magnitudes of less than 15.5. All observations for the pilot study
were carried out on the SOAR 4.1m telescope, and a magnitude limit of 15.5 ensured
36

that we could gather data on the SOAR telescope as efficiently as possible in order
to successfully demonstrate the efficacy of the new target-selection method.
The GPE indices of the HES spectra of known CEMP stars from Aoki et al. (2007)
were calculated in order to estimate the minimum value of the GPE index which might
be representative of a CEMP star. The GPE indices for all of the stars which met the
first set of criteria in the paragraph above are plotted against the (J−K)0 colors in
Figure 3.2. Here, the (J−K)0 color is used as a replacement for temperature, which
is a key determinant of the strength of molecular carbon features. Also shown in the
figure are the known CEMP stars from Aoki et al. (2007). Based on the values of
the GPE indices for the Aoki et al. (2007) stars, a minimum GPE value of 30 ˚ was
A
adopted for target selection.
120

Aoki et al. 2007

100

GPE (Å)

80

60

40

20

0

−20
0.2

0.3

0.4

0.5
(J−K)0

0.6

0.7

0.8

0.9

Figure 3.2: GPE index versus (J−K)0 color for the selected targets (gray dots). The
black dots are the known CEMP stars from Aoki et al. (2007). This figure is from
Placco et al. (2010).
After these preliminary selection criteria were implemented, we were left with 6018

37

candidate CEMP stars. Visual inspection of these candidates was performed in order
to classify the candidate stars into categories. Table 3.1, taken from Placco et al.
(2010), lists the number of candidates for each different classification category.

Table 3.1: Classification of CEMP Candidates

Tag

Description

Number of Candidates

mpca

Absent Ca II K line

4

mpcb

Weak Ca II K line

280

mpcc

Strong Ca II K line

4614

unid

Ca II K line not found

143

fhlc

Faint high-latitude carbon stars

30

habs

Strong absorption H lines

73

hbab

Horizontal-branch/A type star

218

nois

Low S/N

277

ovl

Overlapping spectra

79

art

Artifacts on photographic plates

123

The GPE and (J−K)0 colors of the different types of candidates are plotted in
Figure 3.3. The bluest stars (at lower values of (J−K)0 ) have very strong hydrogen
lines and have likely entered the pool of candidates due to a strong Hγ line at 4340
˚, which is within the wavelength range across which the GPE index is calculated.
A
We do not expect artificial enhancement of the GPE index at higher (J−K)0 , as the
strength of the H lines decreases with temperature. The unid stars tend to be the
reddest colors. At the lower temperatures, there is much less signal in the blue end of
the spectra where the Ca II K line lies, therefore it can be difficult to identify. Also
plotted in Figure 3.3 are the known CEMP stars from Aoki et al. (2007).
Many of these candidate CEMP stars had very strong Ca II K lines, which may
38

120
110
100

mpcc
mpcb
unid
fhlc
habs+hbab
mpca
Aoki et al. 2007

GPE (Å)

90
80
70
60
50
40
30
0.2

0.3

0.4

0.5
(J−K)0

0.6

0.7

0.8

0.9

Figure 3.3: The distribution of the GPE indices versus (J−K)0 color for the CEMP
candidates, according to the classifications defined in Table 3.1. This figure is from
Placco et al. (2010).
have proven to be more metal rich than our target candidates. In order to eliminate
metal-rich candidate stars, we considered the relationship between the KP index and
the BVHES or (J−K)0 parameter space demonstrated by Christlieb et al. (2008).
We selected a cutoff of [Fe/H] = −2.0, as the errors in the measurement of the KP
index (Beers et al., 1999) ensure that we will still have stars enter the sample that
are between [Fe/H] of −1.0 and −2.0. The top two panels of Figure 3.4 show the
relationship between KP and color for constant values of metallicity. The bottom two
panels illustrate which mpcc candidates are removed from consideration due to the
adopted metallicity cutoff. The mpcc candidates that were still included as targets
for this survey are the ones which fall below the metallicity cutoff line for at least one
of the index-color relationships. Once this final selection was achieved, we were left
with 669 CEMP candidates.

39

Figure 3.4: Top panels: Fits of constant metallicity ([Fe/H] = −2.0 and −2.5) in
index-color plots from Christlieb et al. (2008). Bottom panels: Candidate CEMP
stars plotted to show the cutoff region at [Fe/H] < −2.0. The black dots are the
mpca, mpcb, and unid candidates. The gray dots are the mpcc candidates. This
figure is from Placco et al. (2010).

3.4

Observations and Data Reduction

After the final selection of candidates, we observed a large number of the 669 candidate CEMP stars in order to determine [Fe/H] and [C/Fe] so that we could evaluate
the efficiency of this new technique. Observations of the first candidates were obtained
from 2008 to 2010. We employed the Goodman High-Throughput Spectrograph on
the SOAR 4.1 m telescope for these preliminary studies. With a 600 l mm−1 grating
and a 1.03” slit, we obtained medium-resolution optical spectra covering the wavelength range 3600 − 6000 ˚. Typical exposure times ranged from 300 − 1800 s.
A
Each night, the calibrations included 10 bias frames, 10 quartz flats, and HgAr

40

and Cu arc lamps which were taken after every target observation. We also observed
several radial velocity standard stars each night. The data were reduced with standard
IRAF packages, and reduction steps included bias subtraction, overscan correction,
flat-fielding, bad-pixel correction, wavelength calibration, and spectral extraction.
Radial velocities were also determined for each of the objects. The (J−K)0 color,
GPE, and category of each star is listed in Tables 3.2-3.7. Selected reduced Goodman
spectra of these stars are shown in Figures 3.5, 3.6, and 3.7.

Table 3.2: Color, GPE, and Category of Pilot Stars I

Star

(J−K)

˚
GPE(A)

Category

HE 0008+0049

0.58

32.0

mpcb

HE 0024−0550

0.42

30.4

mpcc

HE 0034−0011

0.36

29.7

mpcc

HE 0035−5803

0.36

31.3

mpcc

HE 0053−0356

0.38

36.5

mpcb

HE 0058+0141

0.26

28.6

mpcb

HE 0100−4957

0.58

37.8

mpcc

HE 0102−0004

0.32

29.3

mpcb

HE 0118−4834

0.37

37.4

mpcb

HE 0156−5608

0.49

30.7

mpcc

HE 0159−5216

0.49

32.6

mpcc

HE 0214−0818

0.31

32.9

mpcb

HE 0307−5339

0.44

43.5

mpcb

HE 0316−2903

0.47

37.8

mpcc

HE 0320−1242

0.42

38.1

mpcb

HE 0322−3720

0.62

39.9

mpcb

HE 0336−3948

0.37

30.7

mpcc

HE 0340−3933

0.34

33.0

mpcc

HE 0345+0006

0.53

30.6

mpcc

41

Table 3.3: Color, GPE, and Category of Pilot Stars II

Star

(J−K)

˚
GPE(A)

Category

HE 0405−4411

0.32

31.8

unid

HE 0414−4645

0.34

34.1

mpcc

HE 0440−5525

0.34

32.0

mpcb

HE 0444−3536

0.49

43.2

mpcc

HE 0449−1617

0.42

31.7

mpcb

HE 0451−3127

0.50

30.4

mpcc

HE 0500−5603

0.80

35.6

mpcc

HE 0509−1611

0.52

41.7

mpcc

HE 0511−3411

0.37

33.1

mpcc

HE 0514−5449

0.31

31.0

mpcb

HE 0518−3941

0.18

30.9

mpcc

HE 0535−4842

0.39

30.5

unid

HE 0536−5647

0.49

31.5

mpcb

HE 0537−4849

0.39

30.5

mpcb

HE 0901−0003

0.43

31.1

mpcc

HE 0910−0126

0.26

28.8

mpcb

HE 0912+0200

0.50

45.4

mpcc

HE 0918−0156

0.84

53.2

mpcc

HE 0923−0323

0.39

30.2

mpcc

HE 0928+0059

0.27

30.9

mpcb

HE 0933−0733

0.38

41.5

mpcb

HE 0948+0107

0.50

31.6

mpcb

HE 0948−0234

0.37

34.0

mpcb

HE 0950−0401

0.34

36.3

mpcb

HE 0950−1248

0.38

33.7

mpcc

42

Table 3.4: Color, GPE, and Category of Pilot Stars III

Star

(J−K)

˚
GPE(A)

Category

HE 1001−1621

0.40

34.4

mpcc

HE 1002−1405

0.36

38.8

mpcc

HE 1007−1524

0.36

32.3

mpcc

HE 1009−1613

0.40

39.6

mpcc

HE 1009−1646

0.40

39.2

mpcc

HE 1010−1445

0.56

30.6

mpcc

HE 1022−0730

0.37

30.2

mpcb

HE 1027−1217

0.43

35.2

mpcb

HE 1039−1019

0.40

36.2

mpcb

HE 1045+0226

0.57

53.6

mpcb

HE 1046−1644

0.55

30.2

mpcb

HE 1049−0922

0.58

48.4

unid

HE 1104−0238

0.90

33.8

unid

HE 1110−1625

0.38

33.4

mpcc

HE 1129−1405

0.48

33.0

mpcc

HE 1132−0915

0.39

31.8

mpcc

HE 1133−0802

0.49

40.4

mpcc

HE 1135−0800

0.54

32.7

mpcb

HE 1137−1259

0.58

35.0

mpcb

HE 1142−0637

0.57

34.3

mpcc

HE 1146−1040

0.50

40.9

mpcc

HE 1146−1126

0.58

34.9

mpcc

HE 1147−1057

0.38

33.2

mpcb

HE 1148−1020

0.41

36.2

mpcb

HE 1148−1025

0.42

38.5

mpcb

43

Table 3.5: Color, GPE, and Category of Pilot Stars IV

Star

(J−K)

˚
GPE(A)

Category

HE 1212−1123

0.29

31.5

mpcb

HE 1217−1054

0.55

38.3

mpcc

HE 1217−1633

0.52

56.5

fhlc

HE 1222−1631

0.57

34.1

mpcc

HE 1223−0930

0.50

45.8

fhlc

HE 1224−0723

0.41

38.9

mpcc

HE 1224−1043

0.36

33.5

mpcc

HE 1228−0750

0.30

30.3

mpcb

HE 1228−1438

0.89

42.6

mpcb

HE 1231−3136

0.33

30.3

mpcb

HE 1255−2734

0.43

36.8

mpcc

HE 1301+0014

0.46

32.6

mpcc

HE 1301−1405

0.48

34.3

mpcc

HE 1302−0954

0.49

32.8

mpcb

HE 1311−3002

0.58

34.8

fhlc

HE 1320−1130

0.34

34.4

mpcb

HE 1320−1641

0.87

43.5

mpcc

HE 1321−1652

0.35

43.3

mpcc

HE 1343+0137

0.41

28.3

mpcc

HE 1408−0444

0.22

31.2

mpcb

HE 1409−1134

0.36

36.4

mpcb

HE 1410−0549

0.25

31.0

mpcb

HE 1414−1644

0.47

32.8

mpcc

HE 1418−1634

0.54

30.5

mpcc

HE 1428−0851

0.53

30.5

mpcc

44

Table 3.6: Color, GPE, and Category of Pilot Stars V

Star

(J−K)

˚
GPE(A)

Category

HE 1430−1518

0.79

45.9

unid

HE 1447−1533

0.83

34.5

mpcc

HE 1448−1406

0.37

30.4

mpcb

HE 1458−0923

0.41

47.9

mpcb

HE 1458−1022

0.54

30.2

mpcc

HE 1458−1226

0.47

43.5

mpcc

HE 1504−1534

0.85

38.9

mpcc

HE 1505−0826

0.25

32.2

mpcb

HE 1507−1055

0.80

39.3

mpcc

HE 1507−1104

0.90

46.3

mpcb

HE 1516−0107

0.43

35.1

mpcb

HE 1518−0541

0.54

32.4

mpcb

HE 1527−0740

0.44

37.8

mpcb

HE 1529−0838

0.38

36.4

mpcb

HE 2025−5221

0.39

39.1

mpcc

HE 2052−5610

0.27

39.9

mpcc

HE 2112−5236

0.52

43.5

mpcc

HE 2140−4746

0.36

30.6

mpcc

HE 2151−0332

0.47

40.0

mpcb

HE 2201−1108

0.29

39.3

mpcb

HE 2207−0912

0.41

36.8

mpcc

HE 2209−1212

0.30

39.8

mpcb

45

Table 3.7: Color, GPE, and Category of Pilot Stars VI

Star

(J−K)

GPE(˚)
A

Category

HE 2219−1357

0.20

30.3

mpcb

HE 2231−0710

0.43

57.4

mpcb

HE 2257−5710

0.51

31.6

mpcc

HE 2353−5329

0.29

33.0

mpcc

Figure 3.5: Goodman spectra of pilot-program stars. This figure is from Placco et al.
(2010).

46

Figure 3.6: Goodman spectra of pilot-program stars. This figure is from Placco et al.
(2010).

3.5

Atmospheric Parameters and Carbon Abundances

Atmospheric parameters (Teff , log g, and [Fe/H]) were determined with a new version of the SEGUE Stellar Parameter Pipeline (SSPP; Lee et al., 2008a,b). These
parameters were then used to determine carbon abundances through spectral-fitting
of the CH G band. Synthetic spectra were generated from Kurucz NEWODF models
(Castelli & Kurucz, 2003), with the same line lists that were used in Sivarani et al.
(2006). For the synthetic spectra, Teff ranges from 3500 to 9750 K, log g from 0.0 to
5.0, [Fe/H] from −2.5 to 0.0, and [C/H] from [Fe/H] −0.5 ≤ [C/H] ≤ +0.5.
In order to fit the G band at 4300 ˚, a chi-square minimization was performed
A

47

Figure 3.7: Goodman spectra of pilot-program stars. This figure is from Placco et al.
(2010).
over the wavelength range 4285−4320 ˚. The initial value of [C/H] was set to the
A
value of [Fe/H], which is the solar value of [C/Fe]. From there, carbon abundance
was increased, while keeping the atmospheric parameters fixed, until the minimum
chi-square was reached. Figure 3.8 shows an example of this fitting technique for one
of the pilot-program stars.
In most cases, the error for the [C/Fe] estimates is around 0.15 dex. Tables 3.83.13 contain the radial velocities, atmospheric parameters, and carbon abundances
for all of the pilot-program stars.

48

Figure 3.8: Spectral fitting of the CH G band for one of the pilot-program stars.
The top panel shows the data (black line) as compared to a synthetic spectrum (red
line) of solar carbon abundance, [C/Fe]. The middle panel is a zoomed-in view of
the range over which chi square is minimized for both the solar carbon as well as the
best-fitting synthetic spectrum. The green line here represents the residuals of the fit.
The bottom panel shows the best-fitting synthetic spectrum, with [C/Fe]= +1.32.
Table 3.8: Atmospheric Parameters and Carbon Abundances for Pilot Stars I

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE0008+0049

−14.5

25.6

4862

3.56

−1.27

−0.32

HE0024−0550

81.3

9.6

5539

2.98

−1.78

0.31

HE0034−0011

−159.1

14.7

5935

2.52

−1.74

1.71

HE0035−5803

83.5

21.9

5807

3.75

−0.70

0.48

HE0053−0356

−22.3

23.4

5752

2.21

−1.84

1.85

HE0058+0141

11.8

14.1

6271

4.08

−0.68

0.48

49

Table 3.9: Atmospheric Parameters and Carbon Abundances for Pilot Stars II

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE0100−4957

179.3

15.1

5971

3.54

−1.03

0.28

HE0102−0004

−118.2

1.0

5815

3.45

−2.45

0.64

HE0118−4834

−45.8

6.8

5610

2.27

−2.42

2.14

HE0156−5608

279.1

9.8

5137

4.27

−2.00

0.47

HE0159−5216

62.1

4.5

5118

1.85

−2.03

0.78

HE0214−0818

16.4

13.2

6177

2.85

−1.02

1.32

HE0307−5339

207.1

35.4

5466

2.45

−2.14

1.34

HE0316−2903

259.5

18.7

5795

3.15

−1.56

1.35

HE0320−1242

106.8

11.9

5750

4.18

−0.43

0.18

HE0322−3720

26.4

21.3

4660

4.55

0.00

0.02

HE0336−3948

149.7

3.2

5952

3.86

−0.53

0.28

HE0340−3933

2.6

2.9

5991

4.03

−0.04

0.01

HE0345+0006

20.6

0.4

5199

2.61

−2.74

0.49

HE0405−4411

107.4

17.2

6180

3.55

−1.23

0.85

HE0414−4645

92.6

21.5

5820

3.78

−1.08

0.33

HE0440−5525

97.1

11.5

6362

3.65

−0.52

0.47

HE0444−3536

175.2

24.9

5198

2.42

−1.66

0.96

HE0449−1617

132.2

8.8

5767

4.42

−0.30

−0.14

HE0451−3127

375.6

21.3

5506

1.85

−2.43

1.49

HE0500−5603

142.2

18.6

5809

3.67

−0.56

0.19

HE0509−1611

119.3

15.4

5108

3.05

−0.49

−0.01

HE0511−3411

105.4

26.0

5873

4.04

−0.38

0.16

HE0514−5449

188.8

11.5

6505

4.19

−0.55

0.28

HE0518−3941

60.8

31.6

6562

3.36

−0.56

0.58

HE0535−4842

89.9

16.4

5935

4.31

−0.58

0.25

HE0536−5647

169.3

11.8

5640

4.08

−0.46

−0.16

50

Table 3.10: Atmospheric Parameters and Carbon Abundances for Pilot Stars III

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE0537−4849

94.0

20.7

5083

4.73

−0.38

0.13

HE0901−0003

64.7

10.4

5077

4.14

−0.26

0.01

HE0910−0126

189.4

14.6

6409

3.24

−2.00

0.92

HE0912+0200

101.7

24.8

5004

4.13

−0.06

0.07

HE0918−0156

124.5

11.3

4114

2.34

−1.00

−0.01

HE0923−0323

124.2

13.3

5576

3.67

−0.74

0.28

HE0928+0059

17.4

2.8

6482

3.67

−0.60

0.54

HE0933−0733

60.8

14.5

5445

4.10

−1.04

0.21

HE0948+0107

513.9

6.3

5081

3.88

−2.27

−0.23

HE0948−0234

197.1

18.1

5807

4.07

−1.02

0.41

HE0950−0401

168.4

4.4

5849

3.62

−1.57

1.32

HE0950−1248

92.4

9.1

5555

4.01

−0.03

−0.17

HE1001−1621

32.8

5.5

6113

4.19

−0.46

0.21

HE1002−1405

98.8

10.0

5653

3.77

−0.40

0.28

HE1007−1524

100.9

13.4

5783

3.36

−0.52

0.27

HE1009−1613

81.1

15.4

5765

3.91

0.01

−0.03

HE1009−1646

24.6

25.0

5818

4.08

−0.52

0.15

HE1010−1445

206.2

9.8

4923

3.31

−0.84

0.00

HE1022−0730

87.6

11.3

6140

3.87

−0.64

0.66

HE1027−1217

176.8

25.8

7705

4.32

−0.10

0.81

HE1039−1019

105.1

10.3

5519

4.14

−0.52

0.12

HE1045+0226

273.0

9.5

4939

1.30

−2.75

1.95

HE1046−1644

−85.5

20.4

4915

3.64

−0.15

−0.07

HE1049−0922

24.3

33.8

4410

4.61

−0.26

0.53

HE1104−0238

192.4

15.4

4075

2.08

−1.28

−0.13

51

Table 3.11: Atmospheric Parameters and Carbon Abundances for Pilot Stars IV

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE1110−1625

123.6

8.2

5605

4.51

−0.11

−0.11

HE1129−1405

209.0

32.9

5129

1.69

−2.29

0.65

HE1132−0915

108.8

14.4

7194

4.13

−0.92

0.81

HE1133−0802

−0.2

9.7

5310

3.41

−0.53

0.22

HE1135−0800

189.0

6.2

5175

2.43

−2.17

0.17

HE1137−1259

123.6

15.3

4689

4.49

−0.60

0.23

HE1142−0637

78.3

20.3

5044

3.55

−0.94

−0.29

HE1146−1040

−44.1

38.5

4928

3.70

−0.62

−0.12

HE1146−1126

332.4

9.5

5041

2.07

−2.27

0.38

HE1147−1057

137.9

9.7

5792

4.23

−0.43

0.18

HE1148−1020

262.2

7.1

5994

3.89

−0.80

0.34

HE1148−1025

153.8

25.3

5718

4.22

−0.61

0.14

HE1212−1123

135.8

4.7

6167

3.67

−1.30

0.55

HE1217−1054

65.5

23.5

4715

4.44

−0.34

0.09

HE1217−1633

3.1

264.9

4546

1.32

−2.36

0.61

HE1222−1631

66.9

15.3

5062

2.10

−2.09

0.31

HE1223−0930

198.6

17.7

4940

1.11

−2.52

2.48

HE1224−0723

71.9

6.3

5355

4.34

−0.24

−0.02

HE1224−1043

303.3

13.0

6228

3.44

−1.53

0.56

HE1228−0750

349.7

1.7

5970

2.55

−1.52

−0.15

HE1228−1438

185.4

3.6

4070

2.02

−1.28

−0.13

HE1231−3136

124.0

12.6

6410

3.38

−1.15

0.93

HE1255−2734

6.9

12.2

5446

2.47

−2.32

1.38

HE1301+0014

37.2

18.8

5260

2.85

−2.55

0.23

HE1301−1405

56.8

7.1

6060

3.91

−0.52

0.27

52

Table 3.12: Atmospheric Parameters and Carbon Abundances for Pilot Stars V

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE1302−0954

155.6

0.8

5193

2.13

−2.42

1.05

HE1311−3002

229.8

2.0

4783

1.06

−2.60

0.94

HE1320−1130

222.1

33.9

6149

3.21

−1.85

1.78

HE1320−1641

108.3

0.5

4051

2.25

−0.73

0.98

HE1321−1652

139.7

11.1

5684

2.51

−2.41

2.16

HE1343+0137

97.3

16.8

6273

3.44

−1.49

0.58

HE1408−0444

144.9

6.8

6622

3.66

−2.72

0.97

HE1409−1134

71.1

7.4

5867

4.38

−0.02

−0.02

HE1410−0549

78.1

21.7

7604

4.17

−0.31

1.04

HE1414−1644

91.0

2.6

5176

2.62

−2.43

0.49

HE1418−1634

128.4

8.8

5237

2.18

−2.20

0.45

HE1428−0851

68.3

2.4

5945

2.31

−2.04

0.03

HE1430−1518

407.2

1.6

4399

1.52

−1.64

0.29

HE1447−1533

61.3

18.6

4181

2.44

−0.59

0.01

HE1448−1406

−177.5

12.1

6467

3.09

−1.37

0.15

HE1458−0923

−290.7

8.3

5473

1.76

−2.32

2.51

HE1458−1022

−93.7

16.6

5025

2.07

−2.26

0.42

HE1458−1226

−11.3

26.3

5122

4.55

−0.27

0.22

HE1504−1534

43.3

10.6

4028

2.20

−0.57

0.04

HE1505−0826

78.3

9.1

6514

3.94

−0.53

0.49

HE1507−1055

169.9

9.4

4178

2.18

−1.66

−0.03

HE1507−1104

117.1

3.4

4057

2.26

−0.65

0.15

HE1516−0107

−14

26.7

5276

2.67

−2.11

0.19

HE1518−0541

58.6

8.9

4808

4.98

−0.63

0.31

HE1527−0740

54.8

8.1

6264

2.97

−2.09

1.13

53

Table 3.13: Atmospheric Parameters and Carbon Abundances for Pilot Stars VI

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE1529−0838

−15.2

26.2

5745

4.32

−0.38

0.13

HE2025−5221

209.5

35.3

5619

2.27

−2.32

2.55

HE2052−5610

303.9

0.9

5840

2.32

−2.15

2.46

HE2112−5236

310.5

4.8

4997

1.33

−1.96

1.21

HE2140−4746

81.7

9.2

5958

3.66

−1.16

0.28

HE2151−0332

−95.5

24.1

5259

1.62

−2.51

1.53

HE2201−1108

−82

6.9

6111

3.87

−1.39

0.49

HE2207−0912

−48.8

17.9

5348

2.87

−2.48

0.43

HE2209−1212

139.5

8.2

6194

3.71

−0.18

0.20

HE2219−1357

142.8

18.7

7091

3.80

−0.49

0.75

HE2231−0710

81.9

34.8

5200

2.48

−0.71

−0.04

HE2257−5710

47.1

16.5

5107

2.02

−2.72

0.97

HE2353−5329

117.4

14.1

6013

2.81

−1.85

2.10

3.6

Results

The pilot program carried out with the Goodman HTS on SOAR proved to be very
successful. Out of all stars observed and analyzed, 36% are carbon-enhanced metalpoor stars (CEMP, with [C/Fe] > +1.0 and [Fe/H] < −1.0). For the very metal-poor
stars, with [Fe/H] < −2.0, our selection technique was even more successful, with
42% of stars which are enhanced in carbon. Figure 3.9 shows the [C/Fe] and [Fe/H]
for the pilot program stars.

54

Figure 3.9: [C/Fe] versus [Fe/H] for the pilot sample. The dashed lines are at [C/Fe]
= +1.0 and [Fe/H] = −1.0 in order to clearly mark the CEMP region of parameter
space.

3.7

Conclusions

A new technique has been developed in order to identify new, previously-undiscovered
CEMP stars from the HES prism spectra. By targeting stars simply based on the
strength of the CH G band using the newly-defined GPE index, temperature biases
and metallicity biases are eliminated, allowing for a comprehensive search for new
CEMP stars for a wide range of Teff and [Fe/H] values. The first 120 candidate
CEMP stars have been observed with the Goodman High-Throughput Spectrograph
on the SOAR 4.1-m telescope. The medium-resolution optical spectra obtained during
these observations were reduced and analyzed in order to determine the atmospheric
parameters Teff , log g, and [Fe/H] as well as carbon abundances, [C/Fe]. In this
55

manner, the efficacy of this selection technique was examined, and a large percentage
of the observed candidates did indeed turn out to be carbon-enhanced metal-poor
(CEMP) stars.

56

Chapter 4:
Improving the Selection
Technique and Extending
the Survey with Multiple
Observatories
4.1

Refining the Selection Criteria

Given the success of the pilot program described in the previous chapter and the large
number of candidate CEMP stars remaining from the HES database, the survey has
been continued. After examining the initial results, ways in which the identification
of new CEMP stars could be improved were realized. The majority of the stars which
did not prove to be metal poor were warmer stars. For these stars, the high GPE
index was primarily due to a strong Hγ line rather than a strong CH G band. In
order to correct for this, a new index similar to that of Smith & Norris (1983) and
Morrison et al. (2003) was employed. This new index, EGP, is based on the flux of
the original 200 ˚ wide band relative to flux of a red band, and is formally defined
A
as:
4400
Iλn dλ
dλ.
EGP = −2.5log 4200
4520
Iλm
4425

Here, Iλn and Iλm are the flux of the feature and red band respectively. The red
band was chosen due to its lack of prominent molecular features. As there are large

57

CN bands on the blue side of the feature, the use of a blue band compared to the
G band feature would not be representative of the true relative strength of the line.
Figure 4.1 shows the definition of the EGP index for 6 HES spectra. It is evident
that in cases where the GPE index for two spectra are very similar, the EGP can
vary by a large percentage. In this way, the use of the EGP eliminates the candidates
which would have normally been chosen based solely on the GPE which have strong
Hγ lines that enhance the feature.

Figure 4.1: The new EGP index for 6 HES stars. The green area is the feature band,
and the red area is the red band used for the calculation of EGP.

58

4.2

New Selection

From the HES database, two types of candidates were chosen for the new selection,
those marked stars and bright. The stars were selected according to the following
criteria. Only stars with (B−V)0 (Christlieb et al., 2008) in the range 0.30 ≤(B−V)0 ≤
1.00 were selected. Using J and K magnitudes from 2MASS (Two Micron All Sky
Survey; Skrutskie et al., 2006), only stars within the range 0.20 ≤(J−K)0 ≤ 0.75 were
selected. From this set, only those spectra which had a signal-to-noise ratio of at
least 5 in the region of the Ca II K line were chosen. After these preliminary cuts, the
color-index selection described in the previous chapter was performed, leaving only
those candidate stars which appeared to fall below [Fe/H] = −2.0 for at least one of
the two color/KP index trends. These selections narrowed the stars subsample to
62,311 candidates.
The bright sources from the HES database are those which are close to or above
the saturation limit. The selection of bright candidates was exactly the same as for
the stars, except that there was no (B−V)0 criteria for these, as saturation hinders
the ability to estimate accurate B and V magnitudes (Frebel et al., 2006). With the
other criteria in place, the bright sample was narrowed down to 18,532 candidates.
When including the bright candidates, it was important to take the effects of
saturation into consideration. These effects are evident in Figure 4.2. The distributions should be the same for both samples, but the saturation of the bright sources
leads to systematically lower values of both indices. To correct for this problem, the
bright sources were separated into nine brightness bins. In each of the nine magnitude
ranges, the distributions of both GPE and EGP for the bright sources were shifted
to match the distribution of the unsaturated stars distributions. These nine shifts
for each index were then fit with a polynomial function, so that a global correction
could be applied to all candidates. The final corrected distributions are shown in
Figure 4.3. Here it is evident that the correction applied to the indices resulted in
59

more consistent distributions for both of the subsets.

Figure 4.2: Distribution of GPE and EGP indices for both the stars and the bright
subsets.
After the indices of the bright stars were corrected for saturation effects, the final
parameter cuts were selected according to the distribution of candidate indices as
compared to the indices for the previously-studied CEMP stars of Aoki et al. (2007)
and Placco et al. (2010). This comparison is shown in Figure 4.4. The final candidates
chosen for the continuing survey were those which had GPE > 31 ˚, EGP > −0.56
A
mag, and (J−K)0 < 0.7. The generously high cut in EGP was chosen to ensure that
as many CEMP stars could be discovered as possible, while still maintaining a large
number of candidates. Following all selection criteria, the total sample included 5,454
candidate CEMP stars.

60

Figure 4.3: Distribution of GPE and EGP indices for both the stars and the bright
subsets after the saturation correction.

4.3

Follow-up Observations: GMOS on Gemini

After the initial success of our pilot program, the survey was continued with the
Gemini Observatory. Given the vast numbers of candidate CEMP stars at almost any
RA, the survey is perfect for such queue-based observations. With the Gemini MultiObject Spectrograph (GMOS), candidate stars were observed primarily with GeminiSouth, given the declination of most of our targets. A few targets were observed with
GMOS on Gemini-North. The initial observations were obtained with the B1200
grating and 1” slit. Later observations were gathered with the B600 grating, given
the fact that it is a more commonly-used grating by other queue-based programs. In
this manner, when the weather conditions were ideal for our program, a quick switch
could be made between programs without any overhead due to grating changes. Both

61

Figure 4.4: Selection criteria for GPE and EGP indices. The solid lines mark the
minimum values of GPE and EGP for which final candidates were selected.
the B1200 and B600 are compatible with the further analysis of the spectra.
Each science observation was accompanied by two quartz flat observations (one
before and one after each star) and a CuAr arc lamp for wavelength calibration. As
the Gemini telescopes are 8-m in size (twice that of SOAR), fainter objects (down to
B=16) were gathered while keeping up with the pace of the previous survey effort on
SOAR.
The data were reduced using the Gemini/gmos IRAF extension package. Reduction steps included mosaic procedures to bridge chip gaps, bias subtraction, overscan
correction, flat-fielding, wavelength calibration, and spectral extraction. Radial velocities were also determined. The reduced spectra were then used to estimate atmospheric parameters with the SSPP. Carbon abundances were also determined using
the same procedure decribed in the previous chapter. Tables 4.1-4.3 list the velocities,
atmospheric parameters, and carbon abundances for the Gemini sample.

62

Table 4.1: Atmospheric Parameters and Carbon Abundances for Gemini Stars I

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE 0515−3358

69.8

8.3

5057

1.99

−0.60

0.26

HE 0532−3819

26.3

13.8

5031

1.69

−1.48

0.94

HE 0548−4508

123.2

2.6

5163

1.05

−2.42

2.67

HE 0854+0105

74.6

2.6

4951

2.10

−0.72

−0.03

HE 0901−0003

60.4

7.3

5333

3.91

−0.72

0.20

HE 0911+0011

1.2

8.0

5555

4.04

−0.73

0.05

HE 0919−0049

53.3

15.7

5609

3.88

−0.35

0.10

HE 0923−0016

55.5

10.7

5930

3.49

−1.27

0.98

HE 0930−1047

43.0

7.0

5387

4.14

−1.06

0.15

HE 0932−0838

11.9

7.3

5648

3.11

−0.56

0.35

HE 0942−0446

−12.2

15.9

5920

3.91

−0.08

0.03

HE 0943−0227

73.6

1.7

5833

3.32

−0.37

0.51

HE 0946−0737

4.2

5.9

5434

4.41

−0.67

0.12

HE 0948−0234

57.9

6.7

5698

4.06

−0.52

0.20

HE 0950−0401

63.8

13.0

6410

2.99

−1.95

3.19

HE 0954−0744

−14.6

7.9

5842

3.77

−0.24

−0.02

HE 1002−1405

7.9

2.6

5731

3.20

−0.45

0.54

HE 1006−1237

21.8

10.6

6315

3.58

−0.44

0.49

HE 1007−1343

−58.2

11.3

5841

4.12

−0.24

0.05

HE 1013−1648

−4.7

10.5

5319

3.08

−0.52

0.27

HE 1016−1625

−26.9

13.1

5188

3.32

−0.64

−0.04

HE 1022−0730

35.1

14.8

5809

3.73

−1.35

0.60

HE 1022−1621

32.2

13.2

5328

4.10

−0.60

0.02

HE 1027−1217

69.3

6.7

7505

3.75

−0.61

2.10

HE 1029−1757

−1.5

28.2

5696

1.73

−3.10

2.90

HE 1031−0020

135.2

33.8

5038

1.23

−2.59

1.74

63

Table 4.2: Atmospheric Parameters and Carbon Abundances for Gemini Stars II

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE 1032−1809

8.5

19.5

4505

4.76

−0.80

0.39

HE 1032−2042

111.6

14.8

5666

3.73

−0.35

0.10

HE 1034−1632

149.1

13.6

6512

3.26

−1.33

1.26

HE 1035−1603

−17.9

17.1

4889

4.88

−0.43

0.11

HE 1037−0301

−24.7

6.2

5616

3.99

−0.39

0.05

HE 1040−1957

18.0

0.0

4513

4.43

−0.59

0.91

HE 1042−1107

72.5

6.6

5705

4.45

−0.63

0.01

HE 1043−1516

−32.2

13.8

4512

4.72

−0.81

0.33

HE 1045−1313

57.1

24.1

5248

3.28

−0.38

0.11

HE 1046−1352

71.2

19.4

5234

1.82

−3.57

3.65

HE 1047−1140

−7.9

23.1

4954

4.44

−0.03

−0.08

HE 1053−2017

−137.5

49.0

5065

4.06

−0.43

−0.02

HE 1054−2718

−20.1

5.7

5515

4.21

−0.12

−0.17

HE 1106−0725

189.1

10.7

5657

3.75

−1.32

0.15

HE 1110−2529

−32.7

26.5

4426

4.77

−1.10

0.25

HE 1111−2817

−4.4

1.6

5375

4.05

−0.26

0.01

HE 1111−3026

227.0

28.5

5036

1.23

−2.02

1.04

HE 1134−1731

144.1

7.5

4832

1.74

−0.88

0.08

HE 1142−1058

59.1

14.8

5541

4.08

−0.59

0.30

HE 1144−2555

−4

7.2

5522

3.89

−0.21

−0.08

HE 1146−1040

−32.4

56.8

5449

4.49

−1.48

0.35

HE 1146−1128

−12.3

1.5

5915

3.42

−0.66

0.38

HE 1147−0415

28.7

2.3

5042

1.05

−2.71

2.58

HE 1150−2800

74.2

2.2

5431

1.70

−1.78

1.53

HE 1153−2326

97.8

0.9

6133

2.79

−1.80

1.75

64

Table 4.3: Atmospheric Parameters and Carbon Abundances for Gemini Stars III

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE 1202−2732

−20.2

12.1

4701

4.81

−0.73

1.01

HE 1216−0739

−13.3

0.7

5335

3.77

−0.26

0.01

HE 1233−2435

75.2

12.7

5612

2.37

−1.65

1.15

HE 1254−2320

−11.3

16.2

5647

3.85

−0.19

−0.04

HE 1304−1128

−49.5

5.2

5269

4.24

−0.45

0.04

HE 1328−1740

−18.7

1.0

5267

4.35

0.14

−0.36

HE 1329−2347

−34.7

20.1

4853

3.27

−0.38

0.06

HE 1337−2608

22.1

7.6

5851

3.09

−2.63

2.25

HE 1343−0626

−64.9

9.4

5533

1.51

−2.09

1.65

HE 1348−3057

157.5

25.0

4964

1.08

−1.91

1.16

HE 1350−2734

9.9

22.0

4373

2.45

−0.89

0.03

HE 1401−1236

81.9

4.4

5565

1.71

−2.12

1.87

HE 1514−0943

−43.9

13.5

6845

3.49

−1.46

2.31

With the addition of the Gemini sample, the success of the survey was enhanced.
Figure 4.5 shows the addition of the Gemini results plotted with the original pilot
sample. The success rate of identifying new CEMP stars was increased after these
observations with 45% CEMP stars at [Fe/H] < −1.0, 55% CEMP stars at [Fe/H]
< −2.0, and 100% CEMP stars at [Fe/H] < −3.0.

65

Figure 4.5: [C/Fe] versus [Fe/H] for pilot sample (black dots) + Gemini sample (blue
dots). The dashed lines are at [C/Fe] = +1.0 and [Fe/H] = −1.0 in order to clearly
mark the CEMP region of parameter space.

4.4

Follow-up Observations: Goodman HTS on
SOAR

In 2010 and 2011, more Goodman observations were obtained according to the new
selection techniques. The details of observations, data reduction, and analysis are
the same as those for the pilot sample described in the previous chapter. Tables 4.44.6 list the velocities, atmospheric parameters, and [C/Fe] for the latest Goodman
observations.

66

Table 4.4: Atmospheric Parameters and Carbon Abundances for Goodman Stars I

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE 0002−1037

13.6

28.1

4896

3.09

−3.07

1.12

HE 0004−2546

70.5

5.7

4916

4.01

−0.59

0.18

HE 0020−2549

64.3

8.9

5078

2.74

−0.91

−0.03

HE 0039−2635

−59.7

5.7

4809

1.05

−3.39

1.64

HE 0046−4712

−50.2

11.2

5340

3.62

−1.04

0.63

HE 0055−2507

74.7

6.6

4895

3.29

−0.34

0.09

HE 0059−6540

62.1

23.2

4727

1.01

−3.17

1.20

HE 0113−3806

−47.7

9.3

5887

3.14

−2.08

1.90

HE 0123+0023

−233.4

15.4

6323

2.97

−1.81

2.05

HE 0134−2504

59.4

28.0

5151

2.43

−3.10

1.85

HE 0317−4705

244.7

31.1

4231

3.06

−3.26

0.95

HE 1444−1219

6.7

1.9

5076

2.49

−2.49

0.24

HE 1503−0918

13.6

9.3

5857

3.48

−1.03

0.78

HE 1508−0736

248.4

87.0

5897

4.19

−0.62

−0.26

HE 2006−5334

−121.9

2.7

4106

2.52

−0.89

−0.02

HE 2030−5323

205.2

16.2

5090

1.79

−2.12

0.85

HE 2030−6056

1.7

5.7

4718

3.43

−0.28

0.03

HE 2033−6206

−28.8

13.2

5956

3.35

−1.31

1.09

HE 2043−5525

137.0

6.0

4870

1.34

−2.65

1.13

HE 2056−6128

−152.3

1.7

5205

3.99

−1.32

0.07

HE 2118−5654

68.5

3.1

5261

1.96

−2.71

1.79

HE 2121−5308

−3.7

13.5

4464

1.05

−2.75

0.50

HE 2125−3447

−75.5

6.2

4469

2.97

−0.40

0.15

HE 2134−0637

−76

1.4

6221

2.95

−2.32

2.63

HE 2134−3940

−33.4

24.1

4700

1.05

−3.01

1.24

HE 2135−0759

77.3

7.3

4742

3.57

−0.27

0.02

67

Table 4.5: Atmospheric Parameters and Carbon Abundances for Goodman Stars II

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE 2136−5928

−25.9

11.2

5163

1.15

−1.85

1.13

HE 2138−5620

−23.2

24.9

5435

4.07

−0.64

−0.21

HE 2146−0247

−191

8.6

5516

3.83

−0.31

−0.19

HE 2148−1058

86.6

5.7

4966

2.93

−0.19

−0.23

HE 2151−0643

−1.5

9.5

5855

4.27

−0.69

0.17

HE 2155−3750

116.1

28.0

4858

1.13

−2.85

1.60

HE 2200−1146

55.9

6.4

6033

4.09

−0.37

0.03

HE 2211−1806

110.9

35.3

4782

1.24

−3.38

1.32

HE 2220−2250

42.5

12.3

4798

2.86

−0.07

−0.36

HE 2229−1619

−176.4

22.0

4957

1.58

−2.51

1.38

HE 2324−0424

−81.3

8.8

5261

3.96

−0.37

−0.24

HE 2339−3236

−95.1

5.9

4643

4.80

−0.61

0.91

HE 0546−4421

50.3

13.1

4966

4.87

0.00

−0.17

HE 0927−0035

128.9

6.6

4737

3.09

−0.43

−0.03

HE 0932+0005

49.1

1.4

4817

4.81

−0.86

0.11

HE 0954+0219

298.7

4.0

5210

1.70

−1.87

1.20

HE 1055−2647

111.3

5.9

4627

2.87

−0.68

0.03

HE 1112−1140

69.7

7.3

4779

3.64

−0.48

0.08

HE 1124−2343

111.9

0.6

5026

4.95

−0.68

0.13

HE 1126−1229

129.6

7.4

4694

3.14

−0.80

0.05

HE 1140−2814

131.6

9.2

4651

2.56

−0.68

−0.08

HE 1141−3140

−78.1

3.5

4689

3.78

−0.44

0.19

HE 1315−2807

193.7

9.1

5530

2.46

0.00

−0.07

HE 1327−2116

292.9

15.6

5191

1.56

−2.35

2.07

HE 1336−1832

11.0

16.0

6270

3.38

−0.08

0.63

68

Table 4.6: Atmospheric Parameters and Carbon Abundances for Goodman Stars III

Star

V (km/s)

σ(V) (km/s)

T eff (K)

log g (cgs)

[Fe/H]

[C/Fe]

HE 1342−2731

10.6

1.0

4825

3.22

−0.39

0.09

HE 1350−2422

281.1

11.6

5412

1.57

−1.56

1.81

HE 1428−1950

73.7

6.9

4563

3.75

−2.07

0.49

HE 1501−0858

−189.9

52.0

4901

4.43

−0.17

−0.08

HE 1507−1122

−45

14.8

4404

2.53

−0.92

−0.21

HE 1509−1437

−71.3

14.6

4424

3.06

−0.74

1.00

HE 1516−0903

−219.2

82.4

4898

4.79

−0.19

−0.06

HE 1523−1155

−32.1

24.8

4741

1.36

−2.83

0.94

HE 1937−6314

−86.7

5.1

5845

4.17

−0.78

−0.02

HE 1939−6626

−97.3

4.1

4760

3.72

0.25

−0.42

The final results for the survey are plotted in Figure 4.6. The black dots and blue
dots are the original pilot sample and Gemini samples, respectively. The red dots are
the most recent 2010 and 2011 Goodman observations. For this entire sample of over
250 stars, 50% are CEMP at [Fe/H] < −1.0, 59% CEMP stars at [Fe/H] < −2.0, and
100% CEMP stars at [Fe/H] < −3.0. The first few years of this survey have resulted
in the identification of ∼ 60 new CEMP stars that have been missed by previous
searches.

69

Figure 4.6: [C/Fe] versus [Fe/H] for pilot sample (black dots) + Gemini sample (blue
dots) + additional Goodman sample (red dots). The dashed lines are at [C/Fe]
= +1.0 and [Fe/H] = −1.0 in order to clearly mark the CEMP region of parameter
space.

70

Chapter 5:
Carbon, Nitrogen, Oxygen, and
Barium Abundances with
XSHOOTER on VLT
5.1

Introduction

There are many known carbon-enhanced metal-poor (CEMP) stars which have incomplete abundance analyses. This number continues to grow rapidly as the survey
described in the previous chapters is continued. More complete abundance analyses
are necessary in order to constrain the progenitor types of the various classes of CEMP
stars, determine the scenarios of chemical enrichment of the Galaxy, and constrain
the IMF.
It is difficult to obtain estimates for carbon, nitrogen, oxygen, and neutron-capture
elements (like barium and strontium) in a timely and efficient manner. Mediumresolution optical spectroscopy of CEMP stars can be used to easily determine carbon
from molecular features like the CH G band. Similarly, nitrogen can be estimated
from optical CN bands, but these features are heavily dependent on the carbon abundance, so it is highly preferable to make use of the NH feature at 3360 ˚. Most
A
instruments do not have sufficiently high efficiency in the blue that is necessary for
nitrogen estimates from this NH band. As explained in Chapter 2, optical spectroscopy can be used for oxygen abundance estimates only with high-resolution and
long exposure times. It is more prudent to make use of the near-infrared CO molec-

71

ular features for oxygen abundances when the carbon abundance is known. Barium
and strontium can be estimated easily as long as the resolution of the spectra is
high enough. Many medium-resolution instruments are unsuitable for estimations of
neutron-capture-element abundances.
These difficulties are eliminated when one observes CEMP stars with the new XShooter spectrograph on the Very Large Telescope (VLT). This instrument is actually
three spectrographs in one. There are near-ultraviolet, visible, and near-infrared arms
for this spectrograph. Spectra in all three of these ranges are gathered simultaneously
with just one observation. For this reason, X-Shooter is ideal for the study of CEMP
stars, since spectra of the optical features of carbon and neutron-capture elements,
the near-UV nitrogen feature, and the near-IR oxygen features can be obtained with
one observation per star.

5.2

Selected Targets and Observations

For the initial study of CEMP stars with this unique instrument, targets were selected
that were known to be metal poor, the majority of which were very metal-poor (VMP)
stars with [Fe/H] < −2.0. These targets were selected from the HERES survey
(Barklem et al., 2005) as well as other samples. The HERES sample already contains
estimates for 15-20 elemental abundances per star.
Observations were carried out in 2010 March and 2010 August. During these
observations, spectra of over 25 of the targets were obtained. For the near-UV arm,
the wavelength coverage was approximately 3000−5500 ˚. A 1” slit was used and the
A
resolving power for these spectra was R = 5100. For the visible arm, the wavelength
coverage was approximately 5500−10000 ˚. The 0.9” slit was used and the resolving
A
power for these spectra was R = 8800. For the near-IR arm, the wavelength coverage
was approximately 10000−25000 ˚. The 0.9” slit was used and the resolving power
A
for these spectra was R = 5100.
72

The raw echelle spectra were reduced with the XSHOOTER pipeline, an automated routine which performs all of the necessary data-reduction steps described in
previous chapters as well as merging the orders. The two-dimensional final output
spectra were extracted using IRAF.

5.3

Model Atmospheres, Line Lists, and Synthetic
Spectra

5.3.1

Model Atmospheres

For each of the observed targets, atmospheric parameters Teff , log g, and [Fe/H]
were previously estimated. Based on these estimates, appropriate MARCS model
atmospheres (Gustafsson et al., 2008) were gathered from the MARCS database
(http://marcs.astro.uu.se/). For each star, a set of model atmospheres with standard compositions were gathered that were calculated using the two closest values of
Teff , log g, and [Fe/H] so that the resultant synthetic spectra could be interpolated
to match the specific parameters for each star. The grid of the available MARCS
atmospheres has ranges of Teff in steps of 250 K, log g in steps of 0.5 dex, and [Fe/H]
in steps of 0.5 dex.

5.3.2

Line Lists

Before the synthetic spectra were generated, line lists for all relevant features were
assembled. The line list for the NH feature at 3360 ˚ is the same of that for Johnson
A
et al. (2007). The CH and CN line lists were generated by Bertrand Plez (private
communication). The CO line lists for the near-IR spectra are those described in
Chapter 2. The hyperfine splitting of the Ba II 5853 ˚ line was compiled by Clare
A
Worley (private communication). All other atomic line data for the line lists come
from the Vienna Atomic Line Database (VALD; Piskunov et al., 1995).
73

5.3.3

Synthetic Spectra

Each of the model atmospheres was used in communion with appropriate line lists to
generate a set of synthetic spectra to be used for abundance determinations. All synthetic spectra were generated using the MOOG stellar line analysis program. Given
model atmospheres and line list inputs, MOOG can be used to directly estimate
abundances of input spectra and thus was used to get first-pass abundance estimates
for the X-Shooter program stars using the closest-matching model atmosphere. The
spectral synthesis feature was further used, however, to generate a set of synthetic
spectra which ranged in carbon and nitrogen abundances for estimation of the C and
N abundances. Each model atmosphere was used to generate 5 synthetic spectra,
with [C/Fe] or [N/Fe] having values of −1.0, 0.0, +1.0, +2.0, and +3.0.

5.4

Abundance Determinations

Synthesis of the CH G band was performed for the [C/Fe] abundance determinations
of the X-Shooter sample stars. Synthetic spectra were interpolated such that the
atmospheric parameters matched those specific to each star. This resulted in 5 final
spectra to be used for the analysis, with [C/Fe] ranging from −1.0 to +3.0. These
five were then interpolated to create 400 new model spectra at a separation of 0.01
dex in [C/Fe]. A chi-square minimization was performed among the synthetic spectra
and observed spectra across the region of the G band (approximately 4280 − 4320 ˚).
A
In this manner, the best-fitting [C/Fe] was selected.
Nitrogen abundances were estimated in the same way by performing spectral
synthesis of the NH band at 3360 ˚. Again, after the interpolation of the spectra, 5
A
were used to create a range of [N/Fe] from −1.0 to +3.0. The best-fitting synthetic
spectrum was then selected after chi-square minimization about the NH feature.
For the cool stars with prominent CO bands in the near-IR spectra, the same

74

technique described in Chapter 2 was used for the determination of [O/Fe]. However,
only two of the CO bands were employed, due to the degradation of the quality of nearIR spectra at the reddest end. For stars with measureable Ba II lines, estimates of
barium abundances were gathered using MOOG exclusively. The model atmosphere
with the closest value of Teff , log g, and [Fe/H] was used as an input in combination
with the Ba II 5853 ˚ line list. Figure 5.1 shows the spectral-fitting for all four of
A
these abundances for one of the program stars.

Figure 5.1: C, N, O, and Ba spectral synthesis for one X-Shooter program star.
For all abundances, a minimum error of 0.25 dex was adopted due to the errors
arising from atmospheric parameter estimates. In cases where there was not very
much difference between the synthetic spectra with abundances of ± 0.5 dex from the
best-fitting spectrum, the error could be as high as 0.5 dex. All X-Shooter abundance
analyses can be found in Appendix A. In all of these figures, the red line is the best75

fitting spectrum, the green line has an abundance of 0.5 dex higher than the estimate,
and the blue line has an abundance of 0.5 dex lower than the estimate.

5.5

Results

Given the range of temperatures of the program stars and the range of quality of
spectra, not every program star has all four abundances determined. However, all
stars have newly-estimated species due to this X-Shooter pilot study. The atmospheric
parameters and barium abundances for the X-Shooter sample are given in Tables 5.1
and 5.2. Nitrogen, carbon, and oxygen abundances are given in Tables 5.3 and 5.4.

Table 5.1: Atmospheric Parameters and Barium Abundances for X-Shooter Stars I

Star

T eff (K)

log g (cgs)

[Fe/H]

[Ba/Fe]

σ([Ba/Fe])

HE 0058−3449

4727

1.94

−2.46

...

...

HE 0206−1916

4741

1.50

−2.04

0.30

0.25

HE 0241−3512

4675

1.20

−1.20

0.85

0.25

HE 0400−2030

5085

0.50

−2.44

> 0.25

...

HE 0408−1733

4260

2.50

−2.20

...

...

HE 0414−0343

4712

2.02

−2.45

0.90

0.25

HE 0430−1609

4651

2.24

−2.25

1.10

0.25

HE 0430−4901

5296

3.10

−2.72

...

...

HE 0440−3426

4886

2.24

−2.22

0.85

0.25

HE 0448−4806

5000

1.81

−3.59

2.25

0.25

HE 0516−2515

4153

0.00

−2.50

...

...

HE 1238−0836

3836

1.50

−2.28

...

...

HE 1315−2035

4683

5.00

−2.52

1.60

0.25

HE 1418+0150

4163

2.90

−1.25

...

...

HE 1430−0919

4627

0.90

−2.99

1.60

0.25

HE 1431−0245

4841

4.50

−2.15

0.75

0.25

76

Table 5.2: Atmospheric Parameters and Barium Abundances for X-Shooter Stars II

Star

T eff (K)

log g (cgs)

[Fe/H]

[Ba/Fe]

σ([Ba/Fe])

HE 2138−1616

4910

4.50

−1.23

...

...

HE 2141−1441

4614

3.50

−1.36

...

...

HE 2144−1832

4250

0.00

−1.50

0.85

0.25

HE 2153−2323

4261

0.50

−2.27

1.00

0.25

HE 2155−2043

4778

2.68

−3.44

...

...

HE 2235−5058

5152

1.99

−2.63

1.20

0.25

HE 2250−4229

4784

2.20

−2.95

...

...

HE 2310−4523

4787

2.09

−2.53

...

...

HE 2319−5228

4431

1.71

−3.07

...

...

HE 2357−2718

4710

4.50

−1.50

...

...

HE 2358−4640

5048

2.80

−2.05

−0.40

0.25

Table 5.3: Nitrogen, Carbon, and Oxygen Abundances for X-Shooter Stars I

Star

[N/Fe]

σ([N/Fe])

[C/Fe]

σ([C/Fe])

[O/Fe]

σ([O/Fe])

HE 0058−3449

−0.12

0.25

−0.26

0.25

...

...

HE 0206−1916

0.81

0.25

1.07

0.50

1.23

0.40

HE 0241−3512

0.62

0.50

0.92

0.25

0.47

0.50

HE 0400−2030

> 3.00

...

1.44

0.50

...

...

HE 0408−1733

−0.66

0.50

0.19

0.50

...

...

HE 0414−0343

1.46

0.25

0.92

0.25

...

...

HE 0430−1609

0.11

0.25

1.14

0.25

0.99

0.25

HE 0430−4901

−0.85

0.25

−0.50

0.25

...

...

77

Table 5.4: Nitrogen, Carbon, and Oxygen Abundances for X-Shooter Stars II

Star

[N/Fe]

σ([N/Fe])

[C/Fe]

σ([C/Fe])

[O/Fe]

σ([O/Fe])

HE 0440−3426

1.14

0.25

1.15

0.25

1.82

0.40

HE 0448−4806

1.16

0.25

2.03

0.25

...

...

HE 0516−2515

...

...

...

...

0.27

0.25

HE 1238−0836

−0.49

0.50

0.49

0.25

0.40

0.40

HE 1315−2035

−0.13

0.25

0.87

0.50

...

...

HE 1418+0150

−0.93

0.25

0.55

0.25

...

...

HE 1430−0919

1.80

0.25

1.82

0.25

1.51

0.40

HE 1431−0245

−0.09

0.25

0.68

0.25

...

...

HE 2138−1616

−0.19

0.25

−0.13

0.25

...

...

HE 2141−1441

−0.32

0.50

−0.23

0.25

...

...

HE 2144−1832

0.92

0.50

0.89

0.25

0.96

0.40

HE 2153−2323

1.92

0.50

1.58

0.50

0.60

0.40

HE 2155−2043

−0.78

0.25

0.09

0.25

...

...

HE 2235−5058

0.84

0.25

1.83

0.25

...

...

HE 2250−4229

−0.83

0.25

−0.04

0.25

...

...

HE 2310−4523

−0.59

0.25

−0.28

0.25

...

...

HE 2319−5228

2.79

0.50

0.60

0.25

1.94

0.40

HE 2357−2718

−0.46

0.25

0.25

0.50

...

...

HE 2358−4640

−0.96

0.25

−0.44

0.25

...

...

Scatter plots of all four abundances for all of the program stars are shown in
Figure 5.2. For the following discussion, we adopt the definition of CEMP stars used
in Carollo et al. (2011), [C/Fe] > +0.7 and [Fe/H] < −1.0.
78

Figure 5.2: C, N, O, and Ba abundances for X-Shooter stars.

5.6

Discussion

Of the confirmed carbon-enhanced metal-poor stars from this program, abundance
patterns can be analyzed in order to classify them into subcategories based on neutroncapture-element enhancement. Figure 5.3 shows the abundance patterns for these
CEMP stars.

5.6.1

Neutron-capture-enhanced Stars

All of the CEMP stars plotted in Figure 5.3 appear to be associated with some
form of neutron-capture-element enhancement based on their abundances of Ba. As
discussed earlier, the most prevalent class of CEMP star observed to date is the
CEMP-s category. The CEMP-s stars are believed to be the result of mass transfer

79

Figure 5.3: C, N, O, and Ba abundances for CEMP X-Shooter stars.
from a metal-poor AGB binary companion, and thus it is prudent to compare the
CEMP stars from this study to the abundance patterns of metal-poor AGB stars as
in Chapter 2. Figure 5.4 again shows the CNO abundance patterns for metal-poor
AGB stars of different masses from Herwig (2004).
The stars from the OSIRIS sample discussed in Chapter 2 appeared to be consistent with the lower-mass AGB stars based on their intermediate nitrogen abundances.
Due to hot-bottom-burning (HBB), the higher-mass stars (of 4, 5, and 6 solar masses
in Figure 5.4) produce prodigious amounts of nitrogen at the expense of carbon. Very
few CEMP stars studied to date show the kinds of abundance patterns that would be
consistent with these higher-mass cases. From the X-Shooter sample shown in Figure
5.3, it can be seen that again, most of the CEMP stars are not consistent with the
higher-mass AGB progenitor cases. Rather, they appear to be largely consistent with
80

Figure 5.4: Carbon, Nitrogen, and Oxygen Abundances for AGB stars of different
solar masses (ranging from 2 to 6) from Herwig (2004).
the 2 or 3 solar mass cases, where nitrogen is not enhanced relative to carbon as would
be expected from a HBB scenario. However, there are 2 CEMP stars, HE 0400−2030
and HE 2153−2323 (solid black squares and solid black circles from Figure 5.3, respectively), which do exhibit patterns consistent with a higher-mass AGB companion
progenitor scenario. Note that the Ba abundance for HE 0400−2030 appears too low
to be indicative of a CEMP-s star, but the estimate of [Ba/Fe] > +0.25 is a conservative lower limit, and it is likely that the barium-enhancement is more pronounced.
This particular star has a strikingly high value of [N/Fe] relative to [C/Fe], qualifying
it as a nitrogen-enhanced metal-poor (NEMP) star. Due to the abundance predictions
from AGB models, if metal-poor AGB stars were primarily of the high-mass variety,
these NEMP stars should in fact be more prevalent. However, they are notoriously

81

underabundant within the family of confirmed CEMP-s stars, and indeed previous
efforts to seek out such stars have been difficult (Johnson et al., 2007). The range of
abundance patterns among the CEMP stars in the X-Shooter sample are indicative
of a range of possible progenitor AGB masses.

5.6.2

CEMP-no Stars

All of the CEMP stars in Figure 5.3 have a level of barium-enhancement which
rules out their association with the CEMP-no class. All of these stars have [C/Fe]
abundance estimates of at least +0.7. There is yet one star from the X-Shooter
sample which can be considered a possible CEMP-no star, HE 2319−5228. This
star has [C/Fe] = +0.60 ± 0.25, and therefore meets carbon-enhanced requirements
within errors. The nitrogen and oxygen enhancements are so high ([N/Fe] = +2.79,
[O/Fe]= +1.94) that this star can confidently be considered within the CEMP star
family. This star has no measureable barium in its spectrum, and so it can be classified as a CEMP-no. With [Fe/H] < −3.0, this star is one of the most metal-poor of
the sample, and the lowest-metallicity CEMP stars tend to be of the CEMP-no variety. Even more convincing is the extremely high abundance of nitrogen in this star.
One of the likely progenitor scenarios of CEMP-no stars is the mass loss by rapidly
rotating massive mega metal-poor stars (Hirschi et al., 2006; Meynet et al., 2006).
The theoretical abundance yields of this scenario suggest that efficient CNO cycling
prior to mass loss would produce such extreme enhancements of nitrogen relative to
carbon. This theory has been ruled-out by some previous studies of other individual
CEMP-no stars like that of Ito et al. (2009), where the nitrogen enhancement was too
low relative to carbon to be representative of this scenario. The abundance pattern
of HE 2319−5228, however, appears to be consistent with a progenitor object like the
massive rapidly rotating mega metal-poor stars which might have been present in the
earliest generations of stars.

82

5.7

Conclusions

The X-Shooter spectrograph on VLT is one of the most useful current instruments
for the study of CNO abundances of carbon-enhanced metal-poor stars. The three
spectral ranges available for simultaneous spectroscopy allow for abundance determinations of these elements, as well as Ba, from moderate-resolution spectra alone.
The 27 stars from the initial program were analyzed through spectral synthesis of
molecular CH, NH, and CO bands as well as atomic Ba II lines. With all four of
these abundances known, it is possible to classify each star (e.g. CEMP-s, CEMPno) according to its abundance pattern. Of the metal-poor stars observed for this
program, many of them turned out to be carbon-enhanced. The majority of these
confirmed CEMP stars exhibit enhancements of barium. Comparison of these CEMP
stars to AGB model yields reveal that the progenitor AGB stars were primarily of
the lower-mass variety (as in Kennedy et al. (2011)), while there are 2 CEMP-s stars
which appear to be associated with more massive AGB progenitor types due to their
large abundances of nitrogen relative to carbon. There is one possible CEMP-no star
based on this analysis. The results of these types of studies serve to constrain what
we know about the early formation and subsequent evolution of the Galaxy in both a
physical and chemical sense. The success of this program paves the way for future XShooter programs which will include the observations of the newly-discovered CEMP
stars as a result of the survey discussed in the previous chapters.

83

Chapter 6:
Conclusions
A comprehensive study of Galactic carbon-enhanced metal-poor (CEMP) stars has
been presented in this dissertation. With the use of several different instruments
on several different telescopes, moderate-resolution spectroscopy of near-infrared, optical, and near-ultraviolet features has resulted in accurate estimates of the crucial
CNO abundances.
Near-IR medium-resolution spectroscopic observations with the OSIRIS spectrograph on the SOAR 4.1-m telescope were obtained in order to estimate [O/Fe] for a
sample of 57 CEMP stars. This method of abundance analysis allows for oxygen abundances accurate to about 0.4 dex. Upon comparison of these [O/Fe] estimates to highresolution estimates, it was found that the values of [O/Fe] the sample of 57 CEMP
stars largely fall within regions of parameter space occupied by the high-resolution
estimates of oxygen for other CEMP stars, primarily those which are enhanced in
s-process elements. The [C/Fe], [O/Fe], and [N/Fe] (when available) estimates agree
with the patterns from Herwig (2004) closely enough that mass transfer from an AGB
companion is a likely scenario for many of the stars in our sample, especially when
the effects of dilution are considered.
Large-number statistics are necessary in order to begin to comprehend the complexities of chemical evolution and galaxy formation. To this end, a new technique has
been developed in order to find previously-undiscovered CEMP stars from the HES
prism spectra. The new selection technique is based on a combination of two indices

84

to determine the strength of the CH G band, GPE and EGP. In this manner, temperature biases and metallicity biases are eliminated, and spurious selections due to
prominent hydrogen lines in the spectra are eliminated as well. The candidate CEMP
stars have been observed with the Goodman High-Throughput Spectrograph on the
SOAR 4.1-m telescope as well as the with the GMOS spectrographs on Gemini-South
and Gemini-North 8-m telescopes. The medium-resolution optical spectra obtained
during these observations were reduced and analyzed in order to determine the atmospheric parameters Teff , log g, and [Fe/H] as well as carbon abundances, [C/Fe].
The success rate of this program is remarkably high. Roughly 50% of the candidates
which are metal poor are carbon enhanced, with [C/Fe] > +1.0. About 60% of the
very metal-poor candidates turned out to be enhanced in carbon, and 100% of the
extremely metal-poor candidates are carbon enhanced.
As the survey effort continues, there is a need to gather more complete information for each of the newly-discovered CEMP stars. It has been demonstrated here
that the X-Shooter spectrograph on the Very Large Telescope (VLT) provides for unrivaled abundance analyses due to its extremely broad spectral coverage. To test the
capabilities of this new instrument, 27 metal-poor stars from the initial pilot program
were analyzed through spectral synthesis of molecular CH, NH, and CO bands and
the Ba II lines. Upon classification of the confirmed CEMP stars from the sample, it
was discoved that the vast majority exhibit some neutron-capture-element enhancement, indicating that the likely source of carbon enhancement for these stars was the
result of mass transfer from a binary AGB companion. While most of the CEMP-s
stars in the sample appear to be associated with low-mass AGB progenitors, there
are 2 which have such high enhancements of nitrogen as would be indicative of a
higher-mass companion. One of the stars appears to be a CEMP-no star based on its
extremely low metallicity, its high abundances of carbon, nitrogen, and oxygen, and
its absence of detectable barium lines in the optical spectrum.

85

The results of these studies serve to confirm that there were more than one (and
probably several) sites of carbon production at early times. Each of the different
classes of CEMP stars has a very specific abundance pattern. In order to classify
CEMP stars in an absolute way, estimates of neutron-capture-elements are necessary.
However, even in the absence of these estimates, it is possible to draw conclusions
about the different types of carbon-enhanced stars observed based solely on the knowledge of the metallicities and CNO abundances. In order to fully understand how the
elements of life were introduced at early times, complete CNO abundance analyses
are necessary for a very large number of CEMP stars. As the survey effort described
in Chapters 3 and 4 continues, vast numbers of CEMP stars will be discovered and
analyzed using moderate-resolution spectrographs like OSIRIS and X-Shooter in communion with the abundance analysis techniques described in this dissertation.
As more and more CNO abundances are determined, the initial mass function
(IMF) can be more readily constrained. As was shown in previous chapters, the
abundance patterns of CEMP-s stars are heavily dependent on the mass of the AGB
binary companions which were at one time associated with them. Furthermore, the
identification of more of the elusive CEMP-no stars that tend to populate the outer
halo helps to uncover the properties of the first generations of stars.

86

APPENDICES

87

Appendix A:
XSHOOTER Spectra: [C/Fe],
[N/Fe], [O/Fe], and [Ba/Fe]
Determinations
A.1

[C/Fe]

Figure A.1: Carbon abundances for X-Shooter stars

88

Figure A.2: Carbon abundances for X-Shooter stars
89

Figure A.3: Carbon abundances for X-Shooter stars
90

Figure A.4: Carbon abundances for X-Shooter stars
91

Figure A.5: Carbon abundances for X-Shooter stars
92

Figure A.6: Carbon abundances for X-Shooter stars
93

Figure A.7: Carbon abundances for X-Shooter stars
94

Figure A.8: Carbon abundances for X-Shooter stars
95

Figure A.9: Carbon abundances for X-Shooter stars
96

Figure A.10: Carbon abundances for X-Shooter stars
97

Figure A.11: Carbon abundances for X-Shooter stars
98

Figure A.12: Carbon abundances for X-Shooter stars
99

Figure A.13: Carbon abundances for X-Shooter stars
100

Figure A.14: Carbon abundances for X-Shooter stars

101

A.2

[N/Fe]

Figure A.15: Nitrogen abundances for X-Shooter stars

102

Figure A.16: Nitrogen abundances for X-Shooter stars
103

Figure A.17: Nitrogen abundances for X-Shooter stars
104

Figure A.18: Nitrogen abundances for X-Shooter stars
105

Figure A.19: Nitrogen abundances for X-Shooter stars
106

Figure A.20: Nitrogen abundances for X-Shooter stars
107

Figure A.21: Nitrogen abundances for X-Shooter stars
108

Figure A.22: Nitrogen abundances for X-Shooter stars
109

Figure A.23: Nitrogen abundances for X-Shooter stars
110

Figure A.24: Nitrogen abundances for X-Shooter stars
111

Figure A.25: Nitrogen abundances for X-Shooter stars
112

Figure A.26: Nitrogen abundances for X-Shooter stars
113

Figure A.27: Nitrogen abundances for X-Shooter stars
114

Figure A.28: Nitrogen abundances for X-Shooter stars

115

A.3

[O/Fe]

Figure A.29: Oxygen abundances for X-Shooter stars

116

Figure A.30: Oxygen abundances for X-Shooter stars
117

Figure A.31: Oxygen abundances for X-Shooter stars
118

Figure A.32: Oxygen abundances for X-Shooter stars
119

Figure A.33: Oxygen abundances for X-Shooter stars
120

Figure A.34: Oxygen abundances for X-Shooter stars
121

Figure A.35: Oxygen abundances for X-Shooter stars
122

Figure A.36: Oxygen abundances for X-Shooter stars
123

Figure A.37: Oxygen abundances for X-Shooter stars
124

A.4

[Ba/Fe]

Figure A.38: Barium abundances for X-Shooter stars

125

Figure A.39: Barium abundances for X-Shooter stars
126

Figure A.40: Barium abundances for X-Shooter stars
127

Figure A.41: Barium abundances for X-Shooter stars
128

Figure A.42: Barium abundances for X-Shooter stars
129

Figure A.43: Barium abundances for X-Shooter stars
130

Figure A.44: Barium abundances for X-Shooter stars
131

Figure A.45: Barium abundances for X-Shooter stars

132

Appendix B:
Acronyms
2MASS . . . . . . Two-Micron All Sky Survey
AGB . . . . . . Asymptotic Giant Branch
CEMP . . . . . . Carbon-Enhanced Metal-Poor
GCE . . . . . . Galactic Chemical Evolution
GMOS . . . . . . Gemini Multi-Object Spectrograph
HBB . . . . . . Hot-Bottom Burning
HES . . . . . . Hamburg/ESO Survey
HIF . . . . . . H-ingestion Flashes
IMF . . . . . . Initial Mass Function
ISM . . . . . . Interstellar Medium
NEMP . . . . . . Nitrogen-Enhanced Metal-Poor
OSIRIS . . . . . . Ohio State Infrared Imager/Spectrometer
SOAR . . . . . . Southern Astrophysical Research
SSPP . . . . . . SEGUE Stellar Parameter Pipeline
VALD . . . . . . Vienna Atomic Line Database
VLT . . . . . . Very Large Telescope
VMP . . . . . . Very Metal-Poor

133

BIBLIOGRAPHY

134

REFERENCES

Alonso, A., Arribas, S., Mart´
ınez-Roger, C. 1996, A&A, 313, 873
Aoki, W., Beers, T. C., Christlieb, N., Norris, J. E., Ryan, S. G., & Tsangarides, S.
2007, ApJ, 655, 492
Barklem, P. S., et al. 2005, A&A, 439, 129
Beers, T. C., Preston, G. W., & Schectman, S. A. 1985, AJ, 90, 2089
Beers, T. C., Preston, G. W., & Shectman, S. A. 1992, AJ, 103, 1987
Beers, T. C., Rossi, S., Norris, J. E., Ryan, S. G., & Shefler, T. 1999, AJ, 117, 981
Beers, T. C. & Christlieb, N. 2005, ARA&A, 43, 531
Beers, T. C., et al. 2007a, ApJ, 168, 128
Beers, T. C., Sivarani, T., Marsteller, B., Lee, Y. S., Rossi, S., & Plez, B. 2007b,
ApJS, 133, 1193
Bekki, K., & Chiba, M. 2001, ApJ, 558, 666
Boothroyd, A. I., Sackmann, I. J., & Ahern, S. C. 1993, ApJ, 416, 762
Campbell, S. W. & Lattanzio, J. C. 2008, A&A, 490, 769
Carollo, D. et al. 2007, Nature, 450, 1020
Carollo, D. et al. 2011, arXiv:1103.3067
Castelli, F., & Kurucz, R. L. 2003, Modelling of Stellar Atmospheres, 210, 20P
Cescutti, G., & Chiappini, C. 2010, arXiv:1004.0088v1
Chiba, M., & Beers, T.C. 2001, ApJ, 549, 325
Christlieb, N., Green, P.J., Wisotzki, L., & Reimers, D. 2001, A&A, 375, 366
Christlieb, N. 2003, RvMA, 16, 191
Christlieb, N., Sch¨rck, T., Frebel, A., Beers, T. C., Wisotzki, L., & Reimers, D.
o
2008, A&A, 484, 721
Cohen, J. G. et al. 2005, ApJ, 633, L109
Collet, R., Asplund, M., & Trampedach, R. 2006, ApJ, 644, L121
Collet, R., Asplund, M., & Trampedach, R. 2007, A&A, 469, 687
135

Cooke, R., Pettini, M., Steidel, C. C., Rudie, G. C., & Jorgenson, R. A. 2011, MNRAS, 412, 1047
Cristallo, S., Straniero, O., Gallino, R., Piersanti, L., Dom´
ınguez, I., & Lederer, M.
T. 2009a, ApJ, 696, 797
Cristallo, S., Piersanti, L., Straniero, O., Gallino, R., Dom´
ınguez, I., & K¨ppeler, F.
a
2009b, , 26, 139
Denissenkov, P. A., & VandenBerg, D. A. 2003, ApJ, 593, 509
Denissenkov, P. A. & Pinsonneault, M. 2008, ApJ, 679, 1541
Depoy, D. L., Atwood, B., Byard, P. L., Frogel, J., & O’Brien, T. P. 1993, Society of
Photo-Optical Instrumentation Engineers (SPIE) Conference Series, 1946, 667
Fabbian, D., Nissen, P. E., Asplund, M., Pettini, M., & Akerman, C. 2009, A&A,
500, 1143
Forestini, M. & Charbonnel, C. 1997, A&AS, 123, 241
Frebel, A. et al. 2005, Nature, 434, 871
Frebel, A. et al. 2006, ApJ, 652, 1585
Girardi, L., Bressan, A., Bertelli, G., & Chiosi, C. 2000, A&AS, 141, 371
Goswami, A., Aoki, W., Beers, T. C., Christlieb, N., Norris, J. E., Ryan, S. G., &
Tsangarides, S. 2006, MNRAS, 372, 343
Goswami, A., Karinkuzhi, D., & Shantikumar, N. S. 2010, MNRAS, 402, 1111
Gustafsson, B., Edvardsson, B., Eriksson, K., Jørgensen, U. G., Nordlund, A., &
Plez, B. 2008, A&A, 486, 951
Hajduk, M. et al. 2005, Science, 308, 231
Herwig, F. 2003, ASP Conf. Ser., 304, 318
Herwig, F. 2004, ApJS, 155, 651
Herwig, F. 2005, ARA&A, 43, 435
Herwig, F., Pignatari, M., Woodward, P. R., Porter, D. H., Rockefeller, G., Fryer, C.
L., Bennett, M., & Hirschi, R. 2010, arXiv:1002.2241v1
Hirschi, R., Fr¨hlich, C. Liebend¨rfer, M., & Thielemann, F.-K. 2006, RvMA, 19, 101
o
o
Ito, H., Aoki, W., Honda, S., & Beers, T.C. 2009, ApJ, 698, L37
Johnson, J. A., Herwig, F., Beers, T. C., Christlieb, N. 2007, ApJ, 658, 1203
136

Karakas, A. I. 2010, MNRAS, 403, 1413
Karakas, A. I., Campbell, S. W., & Stancliffe, R. J. 2010, ApJ, 713, 374
Kennedy, C. R., et al. 2011, AJ, 141, 102
Kurucz, R. L. 1993, Kurucz CD-ROM 15, Diatomic Molecular Data for Opacity
Calculations (Cambridge: SAO)
Lau, H. H. B., Stancliffe, R. J., & Tout, C. A. 2009, MNRAS, 396, 1046
Lee, Y. S., et al. 2008, AJ, 136, 205
Lee, Y. S., et al. 2008, AJ, 136, 2022
Lucatello, S., Beers, T. C., Christlieb, N., Barklem, P. S., Rossi, S., Marsteller, B.,
Sivarani, T., Lee, Y. S. 2006, ApJ, 652, L37
Lugaro, M., Herwig, F., Lattanzio, J. C., Gallino, R., & Straniero, O. 2003, ApJ, 586,
1305
Marigo, P, Girardi, L., Chiosi, C., & Wood, P. R. 2001, A&A, 371, 152
Marsteller, B. E. ”The Frequency of Carbon-Enhanced Metal-Poor Stars and the
Origin of Carbon in the Universe”, PhD diss., Michigan State University, 2007
Marsteller, B., Beers, T. C., Sivarani, T., Rossi, S., Placco, V., Knapp, G. R., Johnson, J. A., Lucatello, S. 2009, AJ, 138, 533
Masseron, T. et al. 2006, A&A, 455, 1059
Masseron, T., Johnson, J. A., Plez, B., van Eck, S., Primas, F., Goriely, S., Jorissen,
A. 2010, A&A, 509, A93
Mel´ndez, J. & Barbuy, B. 2002, ApJ, 575, 474
e
Meynet, G. & Maeder, A. 2002, A&A, 390, 561
Meynet, G., Ekstr¨m, S., & Maeder, A. 2006, A&A, 447, 623
o
Meynet, G., et al. 2010, in IAU Symposium 268, Light Elements in the Universe (ASP
Conf. Series), eds. C. Charbonnel, M. Tosi, F. Primas, and C. Chiappini, in press
(arXiv:1001.1864)
Morrison, H. L., et al. 2003, AJ, 125, 2502
Norris, J. E., Christlieb, N., Korn, A. J., Eriksson, K., Bessell, M. S., Beers, T. C.,
Wisotzki, L., Reimers, D. 2007, ApJ, 670, 774
Norris, J. E., Gilmore, G., Wyse, R. F. G., Yong, D., & Frebel, A. 2010a, ApJ, 722,
104
137

Norris, J. E., Wyse, R. F. G., Gilmore, G., Yong, D., Frebel, A., Wilkinson, M. I.,
Belokurov, V., & Zucker, D. B. 2010b, ApJ, 723, 1632
Pettini, M., Ellison, S. L., Bergeron, J., & Petitjean, P. 2002, A&A, 391, 21
Pettini, M., Berkeley, J. Z., Steidel, C. C., Chaffee, F. H. 2008, MNRAS, 385, 2011
Placco, V. M. et al. 2010, AJ, 139, 1051
Plez, B. & Cohen, J. G. 2005, A&A, 434, 1117
Piskunov, N. E., Kupka, F., Ryabchikova, T. A., Weiss, W. W., & Jeffery, C. S. 1995,
A&AS, 112, 525
Rossi, S., Beers, T. C., Sneden, C., Sevastyanenko, T., Rhee, J., Marsteller, B. 2005,
130, 2804
Schuler, S. C., Hatzes, A. P., King, J. R., K¨rster, M., & The, L. -S. 2006, AJ, 131,
u
1057
Siess, L., Livio, M., & Lattanzio, J. 2002, ApJ, 570, 329
Sivarani, T., et al. 2006, A&A, 459, 125
Skrutskie, M. F., et al. 2006, AJ, 131, 1163
Smith, G. H., & Norris, J. 1983, PASP, 95, 635
Sneden, C., Cowan, J. J., & Gallino, R. 2008, ARA&A, 46, 241
Spite, M., et al. 2005, A&A, 430, 655
Stancliffe, R. J., & Glebbeek, E. 2008, MNRAS, 389, 1828
Tominaga, N., Umeda, H., & Nomoto, K. 2007, ApJ, 660, 516
Tumlinson, J. 2006, AJ, 641, 1
Umeda, H., & Nomoto, K. 2003, Nature, 422, 871
Umeda, H., & Nomoto, K. 2005, ApJ, 619, 427
Wasserburg, G. J., Boothroyd, A. I., Sackmann, I.-J. 1995, AJ, 447, L37
Woodward, P., Herwig, F., Porter, D., Fuchs, T., Nowatzki, A., & Pignatari, M. 2008,
American Institute of Physics Conf. Ser., 990, First Stars III, 300
Zoccali, M., Hill, V., Lecureur, A., Barbuy, B., Renzini, A., Minniti, D., G´mez, A.&
o
Ortolani, S. 2008 A&A, 486, 177

138