SPECTROSCOPIC ANALYSES OF NEUTRON CAPTURE ELEMENTS IN OPEN CLUSTERS

by

JULIA E. O’CONNELL

Bachelor of Science, 2011 Tennessee State University Nashville, Tennessee

Submitted to the Graduate Faculty of the College of Science and Engineering Texas Christian University in partial fulfillment of the requirements for the degree of

Doctor of Philosophy

December 2017

ACKNOWLEDGEMENTS

I had the pleasure of having lunch with Caty Pilachowski during an American Astro- nomical Society (AAS) meeting in Long Beach, California. During lunch, Caty asked me what I wanted to do with astronomy. She had a right to be curious. She was the reason I was there. I stumbled over all the obvious answers, “I want to teach; I wanted to do research; I’d like to find a way to share my research with the public.” She sat quietly and smiled. After I had exhausted all of the obvious choices, that quiet smile persisted. Finally, I sighed and admitted, “I don’t know, Caty. I just wanted to study astronomy.” That smile turned into a grin as she said, “We all did, dear.” I felt relieved, but somehow part of a larger community. I didn’t quite understand that sense of a larger community then, but I’m quite certain she understood my inability to understand. The conversation casually turned to small-talk about some sights I might be interested in seeing while in Long Beach, beyond the convention.

It was 2009, the International Year of Astronomy. When I found Caty at the meeting, I told her that I had just talked with someone who knew her. “Imagine,” she said, as she turned to introduce me to Chris Sneden. As a graduate student, Chris had developed a tool for stellar spectroscopists that is still the model software for all stellar spectro- scopists, MOOG. Chris asked me what I thought “MOOG” stood for, adding that many had tried to guess what it meant as an acronym. I told him that, in my mind, “MOOG” was a stellar analog for a musical synthesizer. He smiled, and asked to see my work. Since that time, Chris Sneden and Caty Pilachowski have continued to support my work and have, always graciously, offered commentary and advise on analysis techniques, pub- lications and my continuing education. I am deeply grateful that they have been, and will continue to be, a central part of my larger professional community.

In the beginning there was not light, only the realization that I had to earn a physics degree in order to study astronomy. That should have been enough to stop any right- thinking individual, but I told myself if I was terrible at the math I would do something else. As badly as I wanted it, sometimes wanting something just isn’t enough— so much to learn, so much to re-learn. By sheer chance, Stefan Forcey became my Calculus pro- fessor. He made me a better physics student, instantly, when he told me that if two different equations equal the same thing, you can use them anywhere you want. I hadn’t even finished my question. I can’t say if he remembers the conversation, but it was a pivot point for me. Physics was suddenly doable. By design, Stefan became my calculus teacher for the next three years.

As my astronomy educator, Geoff Burks would oscillate between dressing me down in a lecture hall of 70 for missing a class and finding a timely scholarship, REU program or mentor that would keep me funded, interested and working. At one point, so I would have a place to work, he moved a corner bookcase from his home into his tiny office in

ii the Center of Excellence on Tennessee State University campus, and then made sure I had access to the office when he was busy elsewhere. Geoff taught me that I could never imagine the scale of the Universe, but that I could imagine my place in it. Then he made me look. Geoff expected, no, demanded more from me, even as he assured me that I couldn’t know everything. Like any good astronomer, Geoff is always focused on the bigger picture.

As an undergraduate, Matthew Muterspaugh always had something for me to do: Make a graphic for the light path of this interferometer; find the riddle in this article I’m submitting on Friday; make a model of the new guiding head we’re installing on the 2 meter robotic telescope. “We’re installing.” He literally turned up his nose at a report I wrote in Word about the installation. Matthew made me re-write the report in latex typesetting, which I first had to learn. He sent me templates, and waited patiently. In the meantime, we talked. Each conversation a continuation of the last, a progression to- ward being a professional astronomer. Sometimes we talked about what the implications would be if this or that were confirmed or denied. When it came to the topic of life on other planets, he said, “The implications would be profound in either case.”

Matthew introduced me to the gears of science, but Caty introduced me to the sci- ence. The science of stellar spectroscopy. Geoff had suggested I email Caty to start a conversation, one that might lead to a summer REU with her. Of all things, I talked about rainbows. I had no way of knowing, but Caty loves rainbows. That’s how it started, with what was possible, with what light could tell us. We published a paper 2 years later, and I presented preliminary results at the AAS meeting in Long Beach where we had lunch. I am still discovering the depths of her knowledge and dedication to her craft.

I started practicing analysis techniques that Caty had introduced me to at the corner bookcase in Geoff’s tiny office, while Matthew continued to find projects for me, and Ste- fan taught me the unspoken rules of physics. Their dedication is a tribute to educators everywhere, and my gratitude cannot be fully expressed here. Words can be such paltry things at times. I’m a lucky dog.

Peter Frinchaboy eventually became my adviser at Texas Christian University after asking Caty to suggest a graduate student. I say eventually because I resisted coming to a university in Texas, especially one with Christian in its name. Everyone makes mistakes. Peter has guided my graduate student career with both humor and honesty. There’s always a nugget of truth in his humor, and always the flavor of humor in his honesty. I could not have asked for a better combination in an adviser. Thank you, Peter, for your persistence, for your humor, and for your honesty.

As grateful as I am for the guidance and light from my adviser, all the educators and mentors on this journey, I would be sorely remiss if I let this opportunity pass without acknowledging my former graduate student peers. Not only were Sebastian Requena, Pankaj Kumar and Yui Shiozawa my battle comrades in coursework, they challenged

iii me to fill the gaps in their astronomy knowledge, and shared their own application of physics with me. Joseph Campbell said, “If you can see your path laid out in front of you step by step, you know it’s not your path. Your own path you make with every step you take. That’s why it’s your path.” I think this is very true for all graduate students, so it’s entirely reassuring to look over the tall grass to see others, too, struggling to cut their own paths. Battles are won and lost, but all the wins and losses become shared experiences in this process. It is forever my pleasure and honor to know them.

My children, Sally and Danny, have probably sacrificed more than anyone to con- tribute to my journey. Even though both are adults with their own lives, I have missed a good chunk of time with them, their own struggles through adulthood, and missed many holidays and birthdays. They have no regrets, and could not have been more supportive. It speaks to me as much as the title. They call me Dr. Mom. I didn’t think I could love them more. They proved me wrong. Wonderful.

Finally, I cannot end this without giving credit to the two people in my life who, most likely, first put my feet on this path; my father and my brother. After Sputnik went up, my father built his own telescope, ground his own glass for it, and taught himself about light and stars. When my brother was eleven, my father bought him his own telescope, and the two would set off on clear nights to stargaze. I was allowed to tag along, probably because my brother, Bobby, desperately wanted a brother after the addition of his third sister, me. As the last in a line of sisters, I was his designated brother. Almost as soon as I could walk, Bobby dressed me in boy’s clothes and took me along with him to hang out with friends. I was my brother’s sidekick, and allowed access to his telescope without permission. He never called me little sister. Bobby always treated me as his equal. A gift.

My father always bought me the ‘sciency’ gifts for Christmas; the chemistry set, or the filter kit for polarizing light– something along those lines. He taught me about color and light through photography, and walked me through the science kits to show me... possibilities. He also taught me to play chess, how to anticipate a strategy, when to trust my intuition, and when to reason past that intuition. He never let me win. Winning had to be earned– a hard lesson simply taught. That lesson never left me. Thanks, Dad.

As an electrical engineer, Bobby taught me about computer hardware and program- ming in my nascent computer days, and was visibly excited when I told him I was returning to school. He put into words what he had always intimated about me as a child, that I could do anything I set my mind to. He was right, but with one wonderful caveat. I had lots of help. I was allowed to climb onto the shoulders of giants, to see more clearly into my future through their eyes. I’m mixing my metaphors again. Right, Sebastian?

We lost Bobby in 2014, but I can still hear his words. Sometimes words are not paltry things after all. So, this is for you, dear brother, with all my love and gratitude.

iv I extend gratitude to McDonald Observatory, Matthew Shetrone and the observing support staff for help with observations, and to Chris Sneden for helpful discussions con- cerning analysis. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. This research has also made use of the WEBDA database, operated at the Department of Theoretical Physics and Astrophysics of the Masaryk University. Support from the College of Science and Engineering at Texas Christian University for JEO is gratefully acknowledged. This material is based on work supported by the Na- tional Science Foundation under award AST-1311835 to P.M.F.

Funding for SDSS-III & IV has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy Office of Science. The SDSS-III web site is http://www.sdss3.org. The SDSS web site is www.sdss.org.

SDSS-III & IV is managed by the Astrophysical Research Consortium for the Partic- ipating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, Carnegie Mellon University, University of Florida, the French Participation Group, the German Partici- pation Group, Harvard University, the Instituto de Astrofisica de Canarias, the Michi- gan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, Max Planck In- stitute for Extraterrestrial Physics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Prince- ton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale Uni- versity.

v Contents

List of Common Acronyms and Terms in Astronomy xi

1 Introduction1 1.1 Timescales...... 3 1.2 The Periodic Table Past Iron...... 6

2 Motivation for the Study 10 2.1 Why Study Elemental Abundances?...... 10 2.2 Why Elemental Abundances in Open Clusters?...... 12 2.3 A Homogeneous Abundance Analysis...... 14 2.3.1 The Cluster Sample...... 18

3 The Calibration Clusters 21 3.1 Observations and Reductions...... 22 3.2 Analysis...... 26 3.2.1 Radial Velocities...... 26 3.2.2 Stellar Atmosphere Parameters...... 28 3.2.3 Equivalent Width Analysis and Atomic Data...... 34 3.2.4 The Neutron Capture Elements...... 39 3.2.5 Spectrum Synthesis...... 40 3.3 Results and Discussion...... 46 3.4 Literature Comparisons...... 48 3.4.1 M67...... 48 3.4.2 NGC 188...... 52 3.4.3 NGC 7789...... 55 3.4.4 NGC 2420...... 56 3.4.5 NGC 6819...... 58 3.4.6 NGC 6791...... 60 3.4.7 Sources of Uncertainty...... 64 3.5 SUMMARY...... 66

4 A Large Uniform Cluster Sample 71 4.1 The First OCCAM Neutron-Capture Sample of Open Clusters...... 71 4.1.1 Stellar Parameters...... 71 4.1.2 Berkeley 85...... 75

vi 4.2 Summary...... 77 4.3 Discussion...... 84 4.3.1 Chemo-chronology...... 86 4.3.2 The Problem with Barium...... 87

5 Moving Forward 88

A Cluster Line lists 91

B Abundance Measurements Derived for 20 Open Clusters 109

Vita

Abstract

vii List of Figures

2.1 Literature compilation of open clusters...... 16 2.2 Trend of increasing barium abundance with increasing cluster age.... 17

3.1 Color Magnitude Diagram of M67 with Targets...... 23 3.2 Re-derived optical stellar temperatures compared to APOGEE...... 32 3.3 Re-derived optical surface gravity compared to APOGEE...... 33 3.4 Equivalent width...... 35 3.5 Spectrum synthesis example for La II...... 42 3.6 Spectrum synthesis example for Eu II...... 42 3.7 Spectrum synthesis example for Y II...... 43 3.8 Comparison to APOGEE elements common to this study...... 47 3.9 Synthetic Zirconium spectra for NGC 188 – 3336...... 53 3.10 [Fe/H] cluster-to-cluster comparison to literature...... 62 3.11 ∆[X/Fe] cluster-to-cluster comparison to literature...... 63 3.12 Boxplot for all elements considered in this study...... 67 3.13 Temperature dependence for Zr?...... 68 3.14 Zr and Ba as a function of temperature...... 69 3.15 Zr and Ba as a function of metallicty...... 70

4.1 Boxplot including abundances for this study...... 78 4.2 Zirconium as function of temperature and metallicity for the full sample. 79 4.3 Barium as function of temperature and metallicity for the full sample.. 80 4.4 [Ba/Fe] v. log(Age) for the full cluster sample...... 81 4.5 [Ba/Fe] v. log(Age) compared to D‘Orazi et al. (2009)...... 82 4.6 [Ba/Fe] v. log(Age) with giants only from D‘Orazi et al. (2009)..... 83

viii List of Tables

2.1 Final Cluster Sample...... 20

3.1 Observation Log...... 25 3.2 Final Stellar Parameters for Calibration Clusters...... 30 3.3 Literature Comaprision – M67...... 50 3.4 Measured Abundances – M67...... 51 3.5 Literature Comaprision – N188...... 53 3.6 Measured Abundances – NGC 188...... 54 3.7 Literature Comaprision – NGC 7789...... 55 3.8 Measured Abundances – NGC 7789...... 56 3.9 Literature Comaprision – NGC 2420...... 57 3.10 Measured Abundances – NGC 2420...... 57 3.11 Literature Comaprision – NGC 6819...... 59 3.12 Measured Abundances – NGC 6819...... 59 3.13 Literature Comaprision – NGC 6791...... 60 3.14 Measured Abundances – NGC 6791...... 61 3.15 Abundance Sensitivities...... 64

4.1 Paramter Comparison to DR14...... 72 4.2 Measured Paramters and RVs...... 74 4.3 Measured Abundances – Berkeley 85...... 77

A.1 Line List – M67...... 92 A.2 Line List – NGC 188...... 97 A.3 Line List – NGC 2420...... 100 A.4 Line List – NGC 6819 & 6791...... 103 A.5 Line List – NGC 7789...... 106

B.1 Measured Abundances – ASCC 14 & Col 106...... 110 B.2 Measured Abundances – NGC 1912, Be 53, & FSR 498...... 110 B.3 Measured Abundances – NGC 6705...... 111 B.4 Measured Abundances – NGC 1896...... 111 B.5 Measured Abundances – IC 1369...... 112 B.6 Measured Abundances – King 5 & NGC 2355...... 112 B.7 Measured Abundances – Be 9, Be 19, & Be 31...... 113 B.8 Measured Abundances – Melotte 71 & NGC 1798...... 113

ix B.9 Measured Abundances – NGC 6811 & NGC 7062...... 114 B.10 Measured Abundances – Ruprecht 24...... 114 B.11 Measured Abundances – Berkeley 17...... 115 B.12 Measured Abundances – Berkeley 85...... 115

x List of Common Acronyms and Terms in Astronomy

30 • M – Solar Mass =1.9891 x 10 kg

– The subscript refers to a solar value • Metals

– All elements heavier than hydrogen and helium

• Metallicity

– The mass fraction of metals in a star with respect to hydrogen and helium

• [A/B]

– Standard notation for an abundance ratio in log10 solar units, where NA and NB represent the column density number of atoms for elements A and B

– [A/B] ≡ log(NA/NB)star − log(NA/NB) • [Fe/H]

– Often used as a proxy for metallicity in stellar spectroscopic studies

– [Fe/H] = log(Fe/H)star − log(Fe/H) • Abundances

– Formally, log (A) ≡ log(NA/NH)+12.0 for element A • dex – Decimal exponent (deprecated)

– Unit used for abundance value (e.g. [Fe/H] = 0.06 dex)

• AI

– Neutral element A

• AII

– Singly ionized element A

xi • LTE – Local Thermal Equilibrium

– Statistical thermal equilibrium

• Microturbulent velocity (microturbulence, ξ)

– An additional parameter with the dimension of velocity introduced for stel- lar atmospheres to account for some line broadening effects that cannot be measured directly, e.g. Doppler broadening.

• A˚

– Angstrom; 10−10 m

• λ

– Wavelength, usually in A˚

• Gyr

– Gigayear; 109 years

• Myr

– Megayear; 106 years

• ly

– 1 lightyear = 9.46 x 1012 km

• kpc

– 1000 parsecs – 1 parsec (pc) = 3.262 ly

• Main Sequence star

– Star fusing hydrogen (H) to helium (He) in its core

• RGB – Red Giant Branch star

– Evolved star that is no longer fusing H in its core

• HB – Horizontal Branch star

– Evolved star now fusing He in its core (often called a red clump star in open clusters)

• AGB – Asymptotic Giant Branch star

– Evolved star that has run out of both H and He in its core; now shell burning H and He around an inert carbon-oxygen (C-O) core

xii • r-process

– The rapid neutron capture process where the timescale for β decay is longer than the timescale for neutron capture

• s-process

– The slow neutron capture process where the timescale for neutron capture is longer than the timescale for β decay

• SNe II

– Core collapse supernovae, M > ∼8M • SNe Ia

– Sub-class supernovae that occurs in a binary star system

• SNR

– Signal-to-Noise Ratio; also S/N

• APOGEE

– Apache Point Observatory Galactic Evolution Experiment

• SDSS

– Sloan Digital Sky Survey

• 2MASS

– Two Micron All Sky Survey

• IRAF

– Image Reduction and Analysis Facility

xiii • mv or V – Apparent magnitude (or brightness) of a star in the visual band

• Mv – Absolute magnitude in the visual band, or intrinsic brightness of a star as if it were at a distance of 10 pc (parsecs)

• J, H, Ks – The near-infrared magnitudes taken from 2MASS

• B,V

– The blue and visual apparent magnitudes on the Johnson-Cousins system

• B − V ; J − Ks – Difference between magnitudes, also known as the color index

xiv Chapter 1

Introduction

The application of physics in the domain

of astronomy constitutes a line of

investigation that seems to possess almost

unbounded possibilities. In the stars we

examine matter in quantities and under

conditions unattainable in the laboratory.

The increase in scope is counterbalanced,

however, by a serious limitation the stars

are not accessible to experiment, only to

observation, and there is no very direct

way to establish the validity of laws,

deduced in the laboratory, when they are

extrapolated to stellar conditions.

Cecelia Payne, 1925

1 The genesis and evolution for elements of the Periodic Table through stellar nucle- osynthesis has been a key area of astrophysical research for nearly a century. Cecelia

Payne’s work with stellar spectra in 1925 changed our understanding with respect to the abundances of the elements in stars forever (Payne 1925). Aware of the recent work by physicist Meghnad Saha (Saha 1920), Payne successfully described the temperatures of stellar atmospheres, as well as the mass fractions of hydrogen and helium, for the sun and nearby stars. Henry Norris Russell, who had strongly disagreed with her conclusions, was also doing important work in the field of stellar abundances in 1929 using work first for- mulated by Payne (Russell 1929). While more complex in their data structures, targeting techniques, telescopes and instrumentation, large sky surveys that provide abundances for vast numbers of stars today still rely on the first principles for deriving stellar abun- dances, laid down by groundwork in the early decades of the previous century.

Spectroscopic studies and analyses of stars for abundances are, very generally, con- fronted with two concerns: time scales and metallicities. In astronomy, elements heavier than hydrogen and helium are considered “metals.” Metallicity is the measure of the mass fraction of metals in a star with respect to hydrogen and helium. For abundance studies, it is generally given in terms of the relative amount of iron to hydrogen present, determined by analyzing absorption lines in a stellar spectrum. Common practice is to compare a star’s derived metallicity to the star we know best, the Sun. In stellar spec- troscopy, [Fe/H] is used as a proxy for metallicity, scaled to a standard solar value. For this paper, I adopt the standard notations:

2 [A/B] ≡ log(NA/NB)star − log(NA/NB) , (1.1)

where NA and NB represent the column density number of atoms for elements A and B.

And,

[Fe/H]star = log(Fe/H)star − log(Fe/H) . (1.2)

Elements other than iron, for example calcium (Ca) or magnesium (Mg), are then defined in terms of the measured iron to hydrogen ratio:

[X/Fe] = log(NX/NH)star − log(NX/NH) − [Fe/H]star (1.3)

where X represents the element of interest in a star.

1.1 Timescales

Since the seminal papers by Hoyle Hoyle(1954) and Burbidge, Burbidge, Fowler &

Hoyle Burbidge et al.(1957), it is recognized that different elements are produced under different physical conditions, and by stars of different masses with different evolutionary timescales. Stars of low mass, M . 0.60 M , are long-lived with life spans longer the current age of the Universe. Abundances of elements are locked inside these long-lived stars, steadily burning their fuel for 100’s of Gyrs.

3 The first major evolutionary stage of a low to intermediate mass star (0.6M . M .

8M ) is along the Red Giant Branch (RGB). These stars have exhausted the hydrogen

fuel in their cores, but continue to burn hydrogen in a shell around an inert helium core.

Without radiation pressure supplied by the burning core, gravity compresses the core

and releases energy. This energy release, together with the hydrogen burning shell cause

the cooling star to expand its envelope to many times its original size, becoming much

brighter in the process. The convection zone begins reaching to deeper levels inside the

star and “dredging up” material from the interior, called the first dredge up. At 1 M our own Sun will go through this phase. Stars in this mass range have lifespans from

∼50 Myrs to ∼20 Gyrs, the higher mass stars with the shortest lifespans.

High mass stars, M & 8M , live only for 10’s of Myrs, a blink in the cosmic age of the universe, ending their lives as core collapse, Type II supernovae (SNe II). When iron builds up in the core of a high mass star, the consequences are catastrophic. Stellar fusion of elements heavier than iron is endothermic: it requires energy. As an inert iron core builds up, successive layers of silicon, oxygen, neon and carbon consume the remaining fuel of the star. When all energy generation has stopped, these outer layers suffer from gravitational collapse and, within seconds, rebound off the core causing the star to explode as a supernova. In these violent explosions, all the elements produced in their short lives are cast back into the surrounding interstellar medium (ISM) to be recycled into a new generation of stars.

Type Ia supernovae (SNe Ia) differ from the cataclysmic core collapse (SNe II) of massive stars in that it occurs in a binary system of low to intermediate mass stars. The progenitor is believed to be a white dwarf, a stellar remnant whose electron degenerate

4 core is composed mainly of carbon and oxygen. The white dwarf accretes mass from the

companion star until degeneracy pressure can no longer support the weight of the star,

and runaway fusion reactions create enough energy to unbind the star.

One of the last major evolutionary stages of low to intermediate mass stars (0.6M

. M . 8M ) is along the Asymptotic Giant Branch (AGB). These stars end their lives by simply running out of fuel. Having exhausted their core hydrogen (H) and helium

(He), these stars become thermally unstable and “pulse” between alternating ignition of

H and He shell burning around an inert carbon-oxygen (C-O) core. These pulsations are thought to be responsible for “dredging up” material from the interior to the surface.

This is called the third dredge up (TDU). With each thermal pulse (TP) these stars experience mass loss, coughing out the products of their long lives back into the ISM from which they were made.

Except for the lightest elements created by the Big Bang— hydrogen, helium and a smattering of lithium— the chemical composition of the Galaxy has been governed by the synthesis of heavier elements by generation after generation of stars. Each element is synthesized by distinctive mechanisms, but can be arranged into groups with similar means of production. Light elements with even atomic numbers such as oxygen, silicon and calcium are produced by amassing multiples of 4He nuclei. These are the α-elements, so called since their most abundant isotopes are integer multiples of the helium nucleus

(the α particle). The α-elements are produced mainly by SNe II of massive stars (8M .

M . 15M ), with a small contribution from SNe Ia (Woosley & Weaver 1986; 1995, Hoff- man et al. 1999). SNe Ia occur in white dwarfs, the core remnants of low to intermediate mass stars, and are companions to evolved stars in binary star systems.

5 Formally, elements with Z ≤ 13 subject to hydrogen fusion during periods of quiescent stellar evolution are called proton-capture elements, or light elements. In situ production includes C, N, O as well as F, Na, Mg and Al.

Elements with 21 ≤ Z ≤ 30 belong to the Fe-peak group, with odd-Z elements adjacent to the Fe-peak element group having substantially lower abundances. SNe Ia are thought to be the events for the enrichment of Fe-peak elements in the Galaxy (see Truran et al.

2012, Battistini & Bensby 2015, and references therein) A recent work by Ritter et al.

(2017) successfully modeled the observed Galactic trends in odd-Z element abundances from convection feedback in pre-supernova massive stars, and showed that elements K and Cl, light odd-Z elements, could also be created by the same process.

Each of these different mass stars contributes to the chemical evolution of the Galaxy in different ways, depending on the manner of their deaths. Consequently, the timescales for element production is heavily dependent on the mass of stars, for not all stars can produce identical elements in the course of their lifetime.

1.2 The Periodic Table Past Iron

Stellar fusion provides the necessary outward pressure against gravitational contraction, a state called hydrostatic equilibrium, or gravitational equilibrium. Rather than producing energy, however, fusion of elements heavier than iron requires energy. It is endothermic.

Further, Coulomb barriers increase as proton numbers increase, so charged particle re- actions become unlikely. Elements heavier than iron are produced by neutron capture reaction, where the electrostatic repulsion of the Coulomb force ceases to be a barrier.

6 Elements with Z > 30 are produced by neutron capture (n-capture). Neutrons are cap-

tured onto heavy nuclei “seed” material. Unstable nuclei can then β decay, a neutron

becoming a proton, and producing a new element in the process. The basic mechanisms

for n-capture are the slow n-capture process (s-process) and the rapid n-capture process

(r-process). Heavy element production advances beyond the Fe-peak in this manner.

More than half the elements of the Periodic Table are produced by neutron capture.

The r- and s-process are defined on the timescales for n-capture and β decay. The s- process occurs when the timescale for n-capture is longer than the timescale for β decay.

The reverse is true for the r-process, such that the timescale for β decay is longer than the timescale for n-capture. Stable element isotopes are created by single neutron capture in the s-process, where neutron density, or neutron flux, is low. Alternately, the r-process occurs in high neutron flux events, where the buildup of neutrons happens rapidly before unstable nuclei can β decay. This argues for very different production mechanisms and sites.

The production site considered the “main” component for the majority of s-process elements (e.g., Ba and La) has been successfully modeled for low mass (M . 4M )

thermally pulsing asymptotic giant branch (AGB) stars (see Busso et al. 1999; 2004, for

reviews). The primary neutron source in these low mass stars is provided by the 13C(α, n)16O reaction, arising from alternating ignition of H and He shells surrounding the inert

C-O core. While the H-shell is inactive, its convection envelope deepens enough to come in contact with the He inner shell and 12C rich zone (the He-intershell). Some mixing of protons from the H-shell penetrates the top layers of the He-intershell, creating the so-

7 called 13C pocket through the 12C(p, γ)13N(β+, ν)13C reaction1. 13C releases a neutron

through the capture of an α particle, producing 16O in the process (the 13C(α,n)16O

reaction).

Conditions for r-element processes demand high neutron flux environments, and core

collapse supernovae (SNe II) are plausible production site events (Hoyle 1954, Burbidge

et al. 1957). Yet theoretical SNe II yield models face difficulties in reproducing the

observable r-process abundance distribution (Woosley et al. 1994, Sneden et al. 2008,

and references therein). Fortunately, recent observations of a blue “kilonova” (Metzger

et al. 2010, Kasen et al. 2013, Tanaka & Hotokezaka 2013) from a nearby binary neutron

star merger (GW170817) offers a promising production site for r-process nuclei (Smartt

et al. 2017). For a good review of the binary neutron star merger event, see Metzger

(2017, and references therein).

In short, the ratio of the relative abundances of s-process to r-process elements trace

the respective contributions of AGB stellar winds and, quite possibly, Type II SNe and

binary neutron star mergers. Heavy elements build up over time, from one generation of

stars to the next, but not in a simple linear progression. Different elements are produced

by different mechanisms in stars of different masses at different stages of stellar evolution.

Complicating matters, isotopes of elements can be produced along branching pathways,

making it difficult to disentangle production sites by stellar spectroscopy alone. It is

often the case that theory drives observation, and observation then informs theory. The

science of understanding the mechanisms and production sites of elemental abundances

1Standard notation nuclear fusion reactions. The symbols inside the parentheses indicate a particle absorbed or a decay process first, then the particle or decay produced by the reaction last, e.g. 12C + p → γ + 13N and so on.

8 observed in nature has made great strides since 1925, but difficulties exist in modeling the evolutionary stages of stars. For SNe II especially, many uncertainties remain with models: the explosion mechanism itself; the role neutrinos play; modeling multidimen- sional hydrodynamic instabilities. Just modeling the equation of state for this ultradense matter is difficult. Accurate observational evidence via spectroscopic abundance analyses is critical for modeling the production sites for r- and s-process elements successfully.

In this dissertation I provide a summary of current work in the field, and the mo- tivation for choosing each data set. Each chapter stands on its own: In this chapter

I introduce concepts. Chapter 2 describes science goals and motivation, and gives an account of data obtained. In Chapter 3, I present a detailed analysis of well studied open clusters, the “calibration” cluster sample, and provide comparisons to literature for each cluster. Chapter 4 extends the analysis to a larger sample of less well studied open clusters and a discussion of results. In Chapter 5, I give some context for the future.

9 Chapter 2

Motivation for the Study

The principal value of abundance

determinations from stellar spectra is the

clues they give to the nuclear history of

stellar matter and, more generally, of the

matter in the whole Galaxy.

R. Cayrel, 1966

2.1 Why Study Elemental Abundances?

Apart from the natural curiosity to know what stars are made of, the distribution of chemical elements throughout the Milky Way provides specific information about how it has evolved over time. Abundance determinations of individual chemical elements in stars are commonly used as tracers for the formation and evolution of the Galaxy. The code for their evolutionary stage is written in their chemistry, as well as the environment where those stars were born (Hogg et al. 2016, De Silva et al. 2009, Freeman & Bland-

10 Hawthorn 2002). Additionally, by studying the elemental composition of stars, we begin to understand, and gain evidence for, what types of nuclear reactions occur inside of stars, as well as their internal mixing mechanisms.

Each abundance pattern of Fe-peak, r- and s-process and α-elements represents a dis- tinct timescale component of Galactic evolution. Short lived, high mass stars dominated the early chemical evolution of the Galaxy with the production of α-elements, while in- termediate mass stars with longer lives contribute iron-peak elements some time later.

Long lived, low mass stars have preserved much of their original chemical information, and serve as useful fossils for understanding the history and the evolution of the Galaxy.

Many stellar abundance yield models have been suggested to account for the observed r-process distribution in the Galaxy, including neutron star and black hole mergers (Lat- timer et al. 1977, Rosswog et al. 1999, Thompson et al. 2001, Argast et al. 2004, Tanaka

& Hotokezaka 2013), both high and low mass supernovae (Wheeler et al. 1998a;b, Ning et al. 2007), and neutrino driven winds from supernovae (Woosley et al. 1994). None have been entirely successful in modeling the observed abundance distribution for r-process elements (Arnould et al. 2007, Woosley et al. 2002). Sneden et al.(1996) was able to show that very metal-poor stars, [Fe/H] . 3.0, could be excellent tracers for r-process abundances in the early epochs in the Galaxy, but those stars are rare and generally lie outside the Galactic disk.

Galactic chemical evolution (GCE) and stellar abundance yield models rely heavily on accurate high resolution spectroscopic abundance analyses of all elements for model testing, but particularly for elements heavier than iron. Multiple paths exist for the synthesis of different isotopes for some of these elements, and uncertainties still exist

11 in the treatment of convection boundaries, internal mixing and reaction rates for AGBs.

Spectroscopic studies provide observational constraints to models, but relatively few open cluster abundance studies include neutron capture elements. Further, the studies that include n-capture abundance information tend to provide a limited number of isotopes for comparison, and a large scatter exists from study to study. A large, homogeneously derived abundance study, particularly one that includes a variety of n-capture element abundances, help to constrain model parameters.

2.2 Why Elemental Abundances in Open Clusters?

Galactic star clusters provide rich environments to investigate both the chemical evolution of the Galaxy as well as its dynamical properties and, as such, become key tracers for stellar populations research. Open clusters, because they reside almost exclusively in the plane of the Milky Way, are particularly useful systems for understanding Galactic dynamical and chemical evolutionary processes. Moreover, open clusters are an easily discernible Galactic component for which reliable ages can be determined, making it possible to characterize the Galaxy in terms of an age-metallicity relation, age-velocity dispersion relation, as well as the star formation and chemical enrichment history in all parts of the disk (see e.g. Magrini et al. 2009, and references therein). It also allows for an examination of dynamical and chemical evolutionary patterns that may emerge among cluster member stars (Reddy et al. 2013a).

12 Open clusters have classically been used as tracers of Galactic evolution because all stars within a given cluster share the same age and can provide a direct time-line for studying change. Moreover, the Galactic disk hosts hundreds of young and old clusters with ages ranging from several Myr to over 10 Gyr (see a compilation by Dias et al.

2002). They serve as the “laboratories” of astronomy, each cluster a running experiment where we can observe the effects of composition, age, and environment. In aggregate, they serve as tracers for the evolution of the Galaxy, and of its stars, precisely because they are age datable.

However, the inability to establish a uniform, unbiased tracer sample has been a fundamental shortcoming of previous spectroscopic studies. Large sky surveys like the

Sloan Digital Sky Survey/Apache Point Observatory Galactic Evolution Experiment

(SDSS/APOGEE; Eisenstein et al. 2011, Majewski et al. 2017, Blanton et al. 2017) strive to mitigate this deficiency by providing a homogeneous abundance set across an array of nucleosynthetic groups — light, iron-peak and α-elements — for a large number of stellar types throughout the Milky Way. The SDSS-III/APOGEE survey is a high resolution (R ∼22,500), high signal-to-noise (SNR > 100) near-infrared spectroscopic survey (wavelength range from 1.51 to 1.69 µm) which has collected more than 200,000 stars across the full range of the Galactic bulge, bar, disk, and halo using the Sloan 2.5 meter telescope (Gunn et al. 2006). APOGEE is now extending this survey to the south- ern hemisphere (known as APOGEE-2) using the 2.5 meter du Pont Telescope at Las

Campanas Observatory in Chile, and targeting ∼300,000 stars down to H ∼13, rendering a final sample of more than a half million Milky Way stars.

13 Open clusters have also been targeted with APOGEE and APOGEE-2 yielding the

Open Cluster Chemical Abundance and Mapping (OCCAM) survey (Frinchaboy et al.

2013) to determine an unbiased, homogeneously derived metallicity distribution across the Galaxy. The OCCAM survey seeks insight into Galactic evolution trends with de- tailed analyses of light and iron-peak elements determined from the APOGEE infra-red survey. The r- and s-process elements, however, tend to be elusive in infra-red spectra, but critical to our understanding of Galactic evolution.

2.3 A Homogeneous Abundance Analysis

Figure 2.1 is a box plot illustrating the mean abundance ratios from 20 abundance studies of 37 open clusters showing the median values and spread in [Fe/H], odd-Z, Fe-peak, α and neutron capture (n-capture) elements observed, and presented here without regard for age, distance or location in the disk. Spectroscopic studies of open clusters show a large abundance spread for r- and s-process elements, with a notably large scatter in

Ba II (Bragaglia et al. 2001, Yong et al. 2005, Carretta et al. 2007, D’Orazi et al. 2009,

Reddy et al. 2013b, Jacobson & Friel 2013). The large spread in Ba II and other n- capture elements may arise from a several factors, including, but not limited to, diverse analysis methods, different observational instruments and accompanying resolution limi- tations, random errors in equivalent width (EW) measurements and oscillator strengths, or uncertainties arising from the choice of stellar parameters. However, that the large spread is real must be considered as a plausible explanation. A homogeneously derived abundance analysis of these elements may reveal clues for an answer.

14 Alternately, a noticeably smaller spread appears for Y II, Ce II and Nd II, marked as

pale gray in the figure. Data for these elements were gleaned from a scant 3–5 studies,

with Ce II having only ten data points— the minimum to be considered significant for

the figure. It speaks to the need for more studies of these elements, given the importance

of open clusters as tracers for the chemical evolution of the Galaxy.

In an effort to address this need, I performed an investigation into the r- and s- process elemental abundances for a sample of open clusters as part of an optical follow-up to the SDSS-III/APOGEE-1 survey. I have leveraged the APOGEE-based part of the

OCCAM survey to provide cluster membership, allowing for an efficient optical follow-up of known and likely cluster member stars for abundance analyses. Stars were identified as cluster members by the OCCAM survey, which selects member candidates by radial velocity and metallicity from the observed APOGEE sample. To obtain optical data for n- capture abundance analysis, I conducted a long-term observing campaign spanning three years (2013–2016) using the McDonald Observatory1 Otto Struve 2.1 meter telescope and Sandiford Cass Echelle Spectrograph (SES; McCarthy et al. 1993,R ∼60,000). The

SES provides a wavelength range of ∼1400 A,˚ making it uniquely suited to investigate a number of other important chemical abundances, as well as the n-capture elements. For

this study, I derived abundances for 18 elements covering four nucleosynthetic families—

light, iron-peak, n-capture and α-elements, including seven n-capture elements (Y II, Zr

I & II, Ba II, La II, Ce II, Nd II and Eu II), for APOGEE open clusters.

1This paper includes data taken at The McDonald Observatory of The University of Texas at Austin.

15 Of the neutron capture elements, barium (Ba) especially, is enhanced in disk clusters, and there appears to be a large spread over a number of clusters (see Fig. 2.1). Further, an unanticipated trend of increasing barium with decreasing cluster age has been observed

(Jacobson & Friel 2013, Maiorca et al. 2011, D’Orazi et al. 2009, see Figure 2.2).

1.0

0.8

0.6

0.4

0.2 [X/Fe] 0.0

-0.2

-0.4

-0.6 [Fe/H] O Na Mg Al Si Ca Sc Ti V Mn Ni Y Zr Ba La Ce Nd Eu

Figure 2.1. Box plot illustrating the spread in [Fe/H], with α, odd-Z, Fe-peak, and n-capture elements with respect to Fe for 37 open clusters from 20 studies. The middle line of each box indicates the median abundance value, and the upper and lower box boundaries represent the third and first quartiles (75th and 25th percentile) of the data, respectively. The vertical lines represent the full range of abundance values; the dashed line indicates solar. Data for the figure taken from Gonzalez & Wallerstein(2000); Tautvaiˇsieneet al.(2000); Bragaglia et al.(2001); Tautvaiˇsien˙eet al.(2005); Yong et al.(2005); Carretta et al.(2007); De Silva et al.(2007a); Bragaglia et al.(2008); Friel et al.(2010); Pancino et al.(2010); Maiorca et al.(2011); Villanova et al.(2010); Jacobson et al.(2011); Mikolaitis et al.(2011a); Mikolaitis et al.(2011b); Jacobson & Friel(2013); Reddy et al.(2012); Reddy et al.(2013b); Cunha et al.(2017); Overbeek et al. (2016).

16 An AGB yield model by Maiorca et al.(2012) suggests the observed enhanced dis- tribution of neutron capture elements in open clusters may result from a larger than expected, by a factor of four, neutron source reservoir (the 13C(α,n)16O pocket) in low mass (M < 1.5M ) AGBs. These long lived stars may then deliver higher yields of s-process elements to the ISM from which subsequent stars are formed. While there is evidence for an abundance gradient with respect to distance across the Galaxy (Chiappini et al. 2001, Cunha et al. 2016), no age-metallicity relationship has yet been confirmed in open clusters.

0.80 giants (D’Orazi 2009) dwarfs (D’Orazi 2009) 0.60

0.40

[Ba/Fe] 0.20

0.00

-0.20 7.50 8.00 8.50 9.00 9.50 10.00 log Age (yr) Figure 2.2. Reproduced from a study by D’Orazi et al.(2009) which found a trend of decreasing [Ba/Fe] with increasing cluster age for both giant stars and dwarfs like our Sun.

17 To date, the relationship between Ba abundances and cluster age has been presented with literature compilations from different sources. We will explore this relationship with a large, homogeneous data set.

2.3.1 The Cluster Sample

Our aim is to provide a homogeneous, high resolution spectroscopic study of open clus- ters, both well studied and understudied, for a large number of n-capture elements, in conjunction with a number of significant and interesting elements. The subsequent chap- ter will provide abundances for less well studied open clusters. In this chapter I focus on a subset of APOGEE calibration clusters, and provide literature comparisons of abun- dance studies for each cluster, as well as comparisons of several data releases (DRs) for

APOGEE abundance determinations. Table 2.1 details the cluster sample and presents specifics regarding cluster age, distance and position in the Galaxy.

• I have produced a large, homogeneous spectroscopic abundance analysis of neutron

capture elements for not only well studied open clusters, but more importantly, for

clusters that have never been investigated for these elements.

• By deriving a large array of n-capture elemental abundances (Y II, Zr I &II, Ba II,

La II, Ce II, Nd II and Eu II) for this sample size, this study will provide critical

constraints for Galactic evolution models.

• Large sky surveys depend on by-hand analyses of abundances to provide compar-

isons to stellar parameters derived through their abundance pipeline.

18 2.3.1.1 The OCCAM and SDSS- III/APOGEE Surveys

The Sloan Digital Sky Survey- III/Apache Point Observatory Galactic Evolution Ex- periment (SDSS- III/APOGEE, Majewski et al. 2010) is a narrow-band (1.5–1.7 mi- cron), high resolution (R∼ 22, 500), near-infrared spectroscopic survey of approximately

250,000 stars across the Milky Way. The Open Cluster Chemical Analysis and Mapping

(OCCAM, Frinchaboy et al. 2013) survey leverages SDSS- III/APOGEE to produce a comprehensive, uniform, infrared-based dataset for hundreds of open clusters, and to constrain key Galactic dynamical and chemical parameters from this sample.

All target cluster stars for this study have been selected from the SDSS- III/APOGEE data releases 11 & 12 (Alam et al. 2015) and cluster membership for all stars has been determined by the OCCAM survey. The advantage of leveraging these two surveys is that 1) stellar parameters and light element abundances have been well constrained by the APOGEE Stellar Parameters and Chemical Abundances Pipeline (ASPCAP, Garc´ıa

P´erezet al. 2016), 2) while APOGEE provides new uniform Galactic trends for some elements, I have leveraged this survey to add new Galactic trends for key r- and s-process elements and 3) new, never before studied, cluster members have been determined by the OCCAM survey and analyzed for a large array of elemental abundances.

I obtained optical spectroscopic data for 30 open clusters on the 2.1 meter Otto Struve

Telescope and Sandiford Echelle Spectrograph at McDonald Observatory. The second part of the table includes our calibration clusters, open clusters that have been well studied and, therefore, can be compared to multiple literature sources. Our final cluster sample is presented in Table 2.1.

19 Table 2.1. Final Cluster Samplea

Cluster Name RAb Decb Diameter Distance E(B-V) log(Age) Number (hr:min:sec) (deg:min:sec) (arcmin) (pc) (mag.) (yr) of stars

ASCC 14 05 20 31 +35 13 12 26.4 1100 0.18 8.61 1

ASCC 26 06 50 24 +07 15 00 20.4 800 0.13 8.09 2

Berkeley 9 03 32 39 +52 39 14 6.3 1480 0.79 9.30 1

Berkeley 17 05 20 36 +30 36 00 7.0 2700 0.58 10.00 2

Berkeley 19 05 24 06 +29 36 00 4.0 7870 0.32 9.40 2

Berkeley 31 06 57 36 +08 16 00 5.0 8272 0.080 9.31 1

Berkeley 53 20 56 36 +51 05 00 22.0 3100 1.52 9.09 1

Berkeley 85c 20 18 55 +37 45 33 5.0 1760 0.77 9.00 1

Berkeley 91d 21 10 52 +48 32 12 3.0 4100 1.2 9.35 1

Collinder 106d 06 37 06 +05 57 00 35.0 1000 0.12 9.90 2

FSR 0498 00 29 13 +62 24 57 1.5 1800 0.150 8.55 1

IC 1369 21 12 06 +47 44 00 5.0 2083 0.57 8.64 4

King 5 03 14 45 +52 41 12 14.2 1740 0.70 9.10 2

Melotte 71 07 37 30 -12 04 00 7.0 3154 0.11 8.37 2

NGC 1798 05 11 39 +47 41 30 5.0 3470 0.47 9.25 2

NGC 1817 05 12 15 +16 41 24 16.0 1972 0.33 8.61 2

NGC 1896 05 25 42 +29 19 42 20.0 2600 0.60 9.10 2

NGC 1912 05 28 40 +35 50 54 20.0 1400 0.25 8.50 1

NGC 2240 06 33 10 +35 15 00 11.0 1551 0.15 9.20 1

NGC 2355 07 16 59 +13 45 00 7.0 1949 0.22 8.90 1

NGC 6705 18 51 05 -06 16 12 32.0 1877 0.43 8.40 4

NGC 6811 19 37 17 +46 23 18 14.0 1215 0.16 8.80 2

NGC 7062 21 23 27 +46 22 42 5.0 1910 0.43 8.85 1

Ruprecht 24 07 31 53 -12 45 49 5.0 1983 0.35 7.78 2

Calibration Clusters

NGC 2682 (M67) 08 51 18 +11 48 00 25.0 808 0.03 9.45 10

NGC 188 00 47 28 +85 15 18 17.0 1714 0.04 9.88 3

NGC 2420 07 38 23 +21 34 24 5.0 2480 0.04 9.30 6

NGC 6791 19 20 53 +37 46 18 10.0 5035 0.16 9.92 2

NGC 6819 19 41 16 +40 11 49 13.0 2403 0.16 9.36 3

NGC 7789 23 57 24 +56 42 30 25.0 1795 0.28 9.15 5

a Dias catalog, Dias et al.(2002) b J2000 c Possible moving group d Possible open cluster remnant (Bica et al. 2001)

20 Chapter 3

The Calibration Clusters

In total, for the sample of 714 stars, more

than 300 000 equivalent widths were

measured by (the first author’s right)

hand using the IRAF task SPLOT by

fitting Gaussian profiles to the observed

line profiles.

T. Bensby, 2014

Calibration clusters are well studied open clusters in literature that can serve as comparisons to our work. We chose six open clusters previously selected as calibration clusters for APOGEE by M´esz´aroset al.(2013): NGC 2682 (M67), NGC 188, NGC 2420,

NGC 7789, NGC 6791 and NGC 6819. In this chapter, I review the reduction routines and analysis techniques for deriving stellar abundances for all stars in this study. We also provide some information specific to the n-capture elements selected for analysis,

line selection and use, and discuss sources for uncertainties.

21 3.1 Observations and Reductions

As APOGEE is an H band survey, our target members were limited to H ∼10, or V

∼13.5 for the 2.1 meter telescope, although I had some success with fainter stars in directions less affected by reddening, e.g. NGC 6791, E(B − V ) = 0.09 ± 0.04 (Stetson

00 et al. 2004), on nights with exceptional seeing (.1.0 ). Brighter stars in the optical regime are often too bright for the APOGEE infrared region, and dim stars are simply unobservable with the 2.1 meter telescope. Therefore, some clusters observed have a limited number of members. Likely cluster members were culled from the APOGEE data set by the OCCAM survey using both radial velocity (RV) and metallicity, then checked against literature for any discrepancies1. A sample color-magnitude diagram for open cluster M67 and observed targets is featured in Figure 3.1. Using the magnitudes of each star to plot their colors on the x-axis (bluer color to the left, redder to the right) and intrinsic brightness on the y-axis, it is easy to confirm target stars are in an evolving state. All M67 stars targeted for observations, plotted as red filled circles, lie along the red giant branch (RGB). See Table 3.1 for observation dates, exposure times, measured radial velocities, the final SNR of the spectra and abbreviated APOGEE IDs, which will be used throughout this text, e.g. 6329 is APOGEE star 2M01015206+8506329.

1There seems be some question about membership for NGC 188 – 2M01015206+8506329 as Stetson et al.(2004) gives this star a membership probability of 14%, and Platais et al.(2003) reports only a 6% membership probability based on proper motion. Alternately, Jacobson et al.(2011) includes this star as a possible cluster member, and utilizes this star to determine abundance uncertainties for atmospheric parameters (WEBDA ID N188 – 6712, see their Table 6). We find this star to have a radial velocity reasonably consistent with the cluster average at −44.88 km s−1, and metallicity commensurate with our own and previous literature values for this cluster, so I have included star 6329 as a member of NGC 188.

22 6 M67 targets 7 8 9

s 10 K 11 12 13 14 0.2 0.4 0.6 0.8 1.0 (J-Ks)

Figure 3.1. A sample color-magnitude diagram of M67 and selected targets for observa- tion. Using the magnitudes of each star to plot their colors on the x-axis (bluer color to the left, redder to the right) and intrinsic brightness on the y-axis, it is easy to confirm that target stars are in an evolving state. All M67 stars targeted for observations, plotted as filled red circles, lie along the red giant branch (RGB). An isochrone has been added as a guide for the eye.

Observations for all target stars were obtained at McDonald Observatory with the

Otto Struve 2.1 meter telescope instrumented with the SES (R=λ/∆λ≈60,000). Fea- tureless, bright telluric standards were observed each night to remove internal scattered light from the raw data. Several RV standards were also observed over the course of each run. For stars observed through January 2014, the spectrograph setup was centered near λ6175 A˚ and near λ6375 A˚ for stars observed in 2016. Wavelength coverage spans

∼5400–6760 A˚ for the earlier observations and ∼5650–6860 A˚ for stars observed later

23 in the campaign. Unfortunately, the grating angle setting was miscalculated for stars observed in 2013, specifically, all stars in NGC 7789 and four stars in M67, leaving the

Al doublet λλ26696 and 6698 A˚ unavailable for analysis. In general, though, the SES provides 26 orders of continuous wavelength coverage. Consequently, the first and last orders can be safely eliminated without significant loss of expected coverage. Moreover, trimming the first and last order from the spectrum made it possible to exploit lines for analysis very near the spectrum beginning and end. Still, the SES suffers from some internal reflection, referred to as the “picket fence,” which affects ∼15–20 A˚ in a single order per raw image, the order being dependent on the grating angle setting, which in turn determines the wavelength range. Because the spectrograph rides on the back of the telescope, flexure from its weight at zenith distances & 30 degrees can shift the exact location of the picket fence. As a matter of course, any lines considered for analysis that fall within this range were discarded.

Standard IRAF3 routines were employed for data reduction and processing. Specifi- cally, ccdproc was used to apply the bias level correction, trim the overscan region and apply a normalized flat to object frames and comparison arcs. Due to the long expo- sure times, L.A. Cosmic (van Dokkum 2001) was introduced into the reduction sequence for the detection and removal of cosmic rays without altering the flux level of the bias- subtracted, un-scattered light subtracted data. The IRAF package echelle was used for aperture tracing, scattered light removal, extraction of the one-dimensional spectra, flat-

fielding and wavelength calibration (based on a Th-Ar comparison source). Continuum

2shorthand for “wavelengths” 3IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the Na- tional Science Foundation.

24 fitting was performed using a low order spline function for the extracted spectra, which were then co-added to increase the signal-to-noise ratios (SNRs) of the final spectra. The

SNR of the combined spectra for our targets range from ∼75–130, which were used in both radial velocity measurements and abundance analyses.

Table 3.1. Observation Log

a b APOGEE ID ID RA Dec V H UT Date Exp. SNR Vrad

(J2000) (J200) mag. mag. (sec.) (km s−1)

M67 (NGC 2682)

2M08514234+1150076 0076 08:51:42.34 +11:50:07.5 11.63 9.34 2016 JAN 2×2700 125 34.38

2M08514235+1151230 1230 08:51:42.36 +11:51:22.9 11.33 8.85 2016 JAN 2×2700 110 34.44

2M08512618+1153520 3520 08:51:26.18 +11:53:51.8 10.48 8.11 2013 DEC 2×1800 125 34.72

2M08521856+1144263 4263 08:52:18.58 +11:44:26.3 10.46 8.09 2016 JAN 2×1800 110 35.09

2M08514355+1144264 4264 08:51:43.55 +11:44:26.4 10.76 8.29 2016 JAN 2×1800 100 30.14

2M08515952+1155049 5049 08:51:59.52 +11:55:04.7 10.55 8.08 2013 DEC 2×1800 130 35.64

2M08512156+1146061 6061 08:51:21.56 +11:46:06.1 11.44 9.09 2016 JAN 3×1800 105 33.91

2M08514388+1156425 6425 08:51:43.89 +11:56:42.3 10.58 8.11 2013 OCT 2×1800 95 33.49

2M08512990+1147168 7168 08:51:29.91 +11:47:16.8 9.69 6.68 2014 JAN 1×2700 130 32.29

2M08512377+1149493 9493 08:51:23.77 +11:49:49.2 11.52 9.48 2016 JAN 2×1800 110 35.78

NGC 6819

2M19410622+4010532 0532 19:41:06.22 +40:10:53.3 14.07 11.46 2016 OCT 3×2700 105 2.94

2M19412176+4012111 2111 19:41:21.76 +40:12:11.2 13.19 10.34 2016 OCT 3×1800 100 0.92

2M19412707+4012283 2283 19:41:27.08 +40:12:28.3 12.63 9.69 2016 OCT 3×1800 105 0.82

NGC 188

2M00463920+8523336 3336 00:46:39.21 +85:23:33.8 11.54 8.60 2016 JAN 2×1800 90 −39.90

2M01015206+8506329 6329 01:01:52.09 +85:06:32.7 11.67 8.09 2016 JAN 2×1800 115 −42.88

2M01025280+8517563 7563 01:02:52.58 +85:17:56.1 12.68 9.84 2016 JAN 3×1800 95 −42.14

NGC 7789

2M23580015+5650125 0125 23:58:00.15 +56:50:12.5 11.98 8.52 2013 OCT 2×1800 100 −54.08

2M23571400+5640586 0586 23:57:14.00 +56:40:58.6 12.29 9.03 2013 OCT 3×1800 105 −54.39

2M23573184+5641221 1221 23:58:00.15 +56:50:12.5 11.59 8.21 2013 OCT 2×1800 105 −54.85

2M23572242+5641459 1459 23:57:22.42 +56:41:45.9 12.39 9.09 2013 OCT 3×1800 95 −52.45

2M23580275+5647208 7208 23:58:02.75 +56:47:20.8 12.93 9.84 2013 OCT 3×1800 110 −55.38

NGC 2420

Continued on Next Page. . .

25 Table 3.1 – Continued

a APOGEE ID ID RA Dec VH UT Date Exp. SNR Vrad

(J2000) (J200) mag. mag. (sec.) (km s−1)

2M07382114+2131418 1418 07:38:21.15 +21:31:41.8 12.61 10.52 2016 JAN 3×1800 105 73.78

2M07381507+2134589 4589 07:38:15.06 +21:34:59.1 11.29 7.85 2016 JAN 2×1800 110 76.15

2M07382148+2135050 5050 07:38:21.49 +21:35:05.1 13.10 10.81 2016 JAN 3×1800 115 73.82

2M07382195+2135508 5508 07:38:21.96 +21:35:50.9 12.61 10.35 2016 JAN 3×1800 90 72.51

2M07381549+2138015 8015 07:38:15.49 +21:38:01.6 12.65 10.41 2016 JAN 3×1800 95 74.64

2M07382696+2138244 8244 07:38:26.97 +21:38:24.5 12.40 10.06 2016 JAN 3×1800 120 72.89

NGC 6791

2M19204971+3743426 3426 19:20:49.72 +37:43:42.7 13.96 8.17 2015 OCT 3×2700 75 −47.61

2M19210426+3747187 7187 19:21:04.26 +37:47:18.7 13.85 10.31 2015 OCT 3×2700 90 −46.92

a Abbreviated APOGEE ID to be used throughout this chapter

b V magnitudes provided by Simbad Astronomical Database (Wenger et al. 2000)

3.2 Analysis

3.2.1 Radial Velocities

Spectra from RV standards obtained during observing runs served as templates for use

with the IRAF Fourier cross-correlation task fxcor to measure radial velocities, and

associated uncertainties, for all target stars. The observed radial velocities were corrected

for the Earth’s motion using rvcorrect. Errors were generally on the order of δRV = 0.08–

0.80 km s−1. However, our data set includes several identified spectroscopic binaries.

Multiple observations would provide the inherent RV scatter, although apart from the

purpose of this project. Therefore, I assign a global error of ±1.5 km s−1 in an effort to include this as an external uncertainty. For the sake of relevant comparisons, however, I state the mean internal errors by both the authors and this study when discussing cluster velocity averages and individual star RVs.

26 Both Geller et al.(2015, hereafter G15) and Jacobson et al.(2011, hereafter J11)

list M67 – 4264, which corresponds to M67 2014 for G15 and M67 224 for J11, as a

binary member with RV = +32.8 and +35.6 km s−1, respectively. An average of three

observational “visits” for APOGEE DR14 (Abolfathi et al. 2017) found RV = +27.8 km

s−1 with VSCATTER = 0.08 km s−1 for the same star. Our own RV measurement for

this star, +30.1 km s−1, is consistent with the range of RVs found by these other stud-

ies. Similarly, J11 found RV = +23.9 km s−1 for M67 – 9493 (M67 143 in their study),

with APOGEE DR14 RV = +38.2 km s−1 for the same star, a spectroscopic binary with

VSCATTER = 3.9 km s−1 between three visits. This study’s RV measurement for M67

- 9493 = +35.8 km s−1, which is closer to our mean radial velocity for M67, +33.9 ±

0.6 km s−1, in excellent agreement with G15, RV = +33.6 ± 0.03 km s−1 for this clus-

ter. For the same ten stars in M67, the mean APOGEE DR14 RV = +34.0 ± 0.02 km s−1.

For six stars in NGC 2420, I find the mean RV = +73.9 ± 0.13 km s−1, consistent with

Mermilliod et al.(2008, WIYN Open Cluster Study (WOCS)) for 18 stars, RV = +73.6

± 0.15 km s−1. For the same APOGEE stars, RV = +74.2 ± 0.01 km s−1. Gim et al.

(1998) used 597 observations for 112 stars from 1979–1996 to determine membership for

78 giants in NGC 7789. Of these members, 50 constant velocity stars were selected to

determine the mean cluster RV = −54.9 ± 0.12 km s−1. This study finds, for five cluster

NGC 7789 members, mean RV = −54.3 ± 0.09 km s−1. For the same stars, APOGEE

reports RV = −54.7 ± 0.00 km s−1, in excellent agreement with J11 (−54.7 ± 1.3 km s−1

for 26 evolved stars). Also in good agreement are individual radial velocities for NGC

6791 stars associated with this study. Tofflemire et al.(2014) found RVs = −46.5 ± ∼ 0.7

27 km s−1 for both N6791 - 3426 & 7187, for which I find RV = −47.6 & −46.9 ± 0.09 km s−1, respectively. APOGEE DR14 RV measurements for these stars show a difference of only ± 0.20 km s−1 in a one-to-one comparison with our RVs. For NGC 6819 - 0532, 2111

& 2283 I compare to another WOCS radial velocity study. Hole et al.(2009, hereafter

H09) found a cluster average RV of +2.34 ± 0.05 km s−1 for single cluster members, a category to which all three NGC 6819 stars in this study belong. For 0532, H09 found

RV = +3.08 ± 0.13 km s−1, comparable to our +2.94 ± 0.08 km s−1. Similarly, I find good agreement with our measured RVs for 2111 (+0.92 ± 0.16 km s−1) & 2283 (+0.82

± 0.07 km s−1) with H09 2111 RV = +1.14 ± 0.20 km s−1 and RV = +0.56 ± 0.21 km s−1 for 2283.

3.2.2 Stellar Atmosphere Parameters

For this optical follow-up, I re-derived stellar atmospheres for all stars by adopting

4 APOGEE calibrated stellar parameters for temperature (Teff ) and surface gravity (log g) as initial estimates for model stellar atmospheres. The procedure for creating, testing and calibrating APOGEE stellar parameters is thoroughly discussed in M´esz´aroset al.

(2013). Calibration relations for the stellar parameters, abundances and abundance pre- cision is described by Holtzman et al.(2015), while the The APOGEE Stellar Parameter and Chemical Abundances Pipeline (ASPCAP) software details can be found in Garc´ıa

P´erezet al.(2016).

4In astrophysics, the surface gravity of a star is expressed as log g in cgs units, where the unit of acceleration is cm/s2, in base10 logarithmic scale

28 We initially assumed solar metallicity, where log(Fe) = 7.50 (Asplund et al. 2009), and a microturbulence (ξ) value of 2 km s−1 for all stars. All stellar atmospheres were modeled by interpolating in the ATLAS9 grid5 and opacity distribution functions

(ODFNEW, Castelli & Kurucz 2004) for use with the 2017 version of the local thermal equilibrium (LTE) line analysis code MOOG (Sneden 1973). In an iterative process,

MOOG solves for both the Boltzmann and Saha equations given an opacity distribution for a selected temperature and pressure (surface gravity) of a star. Final Teff ’s were attained by removing trends in Fe I abundances with excitation potential.

Surface gravities were refined by both Fe & Ti ionization equilibrium:

− AI A II + e (3.1) where A II is the singly ionized state of neutral element A I. Ionization equilibrium assumes all Fe I and Fe II, and in this case Ti I and Ti II, abundances are equal. Singly ionized absorption lines form deeper in a star’s atmosphere, where temperatures and pressures are greater than at the surface. For higher gravities, pressure forces the atoms and ions closer, increasing the electron density, and thus, the opacity. For lower gravity, the effect is reversed and opacity is decreased. For this reason, ionized lines are sensitive indicators of pressure changes, directly related to surface gravity. For Ti ionization equilibrium, I accepted values for Ti II abundances within 1σ of Ti I, and were used mainly as a check for iron equilibrium.

5Kurucz model atmospheres can be found at http://kurucz.harvard.edu/grids.html

29 When a good estimation for [Fe/H] became evident, metallicity was incorporated into the opacity distribution functions and distilled by an iterative process of finalizing ξ, log g and Teff . We achieved final values by eliminating any remaining bias in Fe I with reduced equivalent width (EW, [log(EW/λ)]) and residual abundances across the entire spectrum. Elemental abundances were derived when a final solution over all parameters was achieved. Our final model atmospheres, including comparisons to ASPCAP stellar parameters (DR14), can be found in Table 3.2. Additionally, DR14 includes, for the

first time, a data driven model for determining stellar parameters and abundances, The

Cannon (Ness et al. 2015, Casey et al. 2016), as an alternative set of stellar parameters.

The Cannon is a machine learning algorithm which determines abundances and stellar parameters independently of stellar model atmospheres, relying instead on the param- eterization of stellar fluxes for “labels” determined from a stellar spectra training set.

Table 3.2 offers a comprehensive comparison for all re-derived optical stellar parameters with DR14 ASPCAP and The Cannon.

Table 3.2. Final Stellar Parameters for Calibration Clusters

This Study DR14 (ASPCAP) DR14 (The Cannon)

Star Teff log g ξ [Fe/H] Teff log g ξ [Fe/H] Teff log g [Fe/H]

(K) (cgs) (km/s) (K) (cgs) (km/s) (K) (cgs)

0076 4775 3.00 1.40 −0.10±0.06 4826 2.99 1.28 0.09±0.01 4829 2.92 0.15±0.06 1230 4720 2.65 1.37 −0.01±0.05 4770 2.86 1.21 0.07±0.01 4781 2.75 0.12±0.06 3520 4700 2.50 1.60 −0.02±0.08 4799 2.48 1.40 0.07±0.01 4790 2.52 0.08±0.06 4263 4750 2.50 1.53 0.01±0.06 4751 2.44 1.49 0.09±0.01 4751 2.51 0.13±0.06 4264 4670 2.55 1.30 0.07±0.05 4732 2.64 1.48 0.13±0.01 4768 2.66 0.18±0.06 5049 4720 2.40 1.50 −0.02±0.07 4789 2.45 1.51 0.06±0.01 4781 2.47 0.09±0.06 6061 4760 2.83 1.51 −0.02±0.07 4777 2.85 1.35 0.12±0.01 4802 2.82 0.17±0.06 6425 4655 2.50 1.50 0.00±0.09 4796 2.44 1.64 0.07±0.01 4796 2.51 0.12±0.06 Continued on Next Page. . .

30 Table 3.2 – Continued

Star Teff log g ξ [Fe/H] Teff log g ξ [Fe/H] Teff log g [Fe/H]

(K) (cgs) (km/s) (K) (cgs) (km/s) (K) (cgs)

7168 4230 1.66 1.60 −0.08±0.05 4303 1.68 1.47 −0.05±0.01 4346 1.85 0.04±0.06 9493 5160 3.20 1.20 0.02±0.05 5082 3.40 1.04 0.02±0.01 5038 3.30 0.01±0.06 0532 4900 2.90 1.50 −0.04±0.07 4910 2.66 1.41 0.12±0.01 4869 3.00 0.13±0.06 2111 4660 2.45 1.44 −0.02±0.07 4650 2.52 1.43 0.14±0.01 4647 2.56 0.16±0.06 2283 4585 2.40 1.50 0.08±0.07 4638 2.50 1.44 0.09±0.01 4612 2.42 0.12±0.06 3336 4300 2.00 1.59 0.06±0.07 4400 2.11 1.42 0.13±0.01 4407 2.14 0.17±0.06 6329 4080 1.60 1.57 0.02±0.07 4129 1.60 1.50 0.15±0.01 4130 1.55 0.17±0.06 7563 4470 2.40 1.46 0.04±0.07 4512 2.45 1.26 0.12±0.01 4512 2.34 0.16±0.06 0125 4350 2.00 1.57 −0.07±0.06 4452 2.04 1.51 0.02±0.01 4463 2.06 0.04±0.06 0586 4520 2.25 1.56 −0.06±0.05 4549 2.30 1.48 0.07±0.01 4545 2.22 0.11±0.06 1221 4365 2.00 1.71 0.02±0.06 4436 2.01 1.55 0.05±0.01 4436 1.98 0.10±0.06 1459 4500 2.26 1.51 −0.02±0.06 4547 2.24 1.49 0.06±0.01 4546 2.20 0.09±0.06 7208 4705 2.40 1.44 −0.03±0.05 4763 2.47 1.43 0.04±0.01 4783 2.55 0.06±0.06 1418 5120 2.90 1.39 −0.18±0.07 5147 2.76 1.43 −0.09±0.01 5037 2.74 −0.16±0.06 4589 4150 1.28 1.63 −0.29±0.05 4183 1.49 1.63 −0.14±0.01 4207 1.54 −0.15±0.06 5050 4885 2.65 1.20 −0.15±0.07 4941 2.59 1.53 −0.10±0.01 4930 2.78 −0.10±0.06 5508 4905 2.40 1.35 −0.09±0.06 4931 2.54 1.47 −0.09±0.01 4994 2.64 −0.05±0.06 8015 4860 2.40 1.48 −0.23±0.05 4922 2.52 1.44 −0.14±0.01 4978 2.65 −0.12±0.06 8244 4850 2.35 1.49 −0.16±0.05 4884 2.48 1.56 −0.11±0.01 4908 2.54 −0.09±0.06 3426 3600 1.00 1.40 0.36±0.04 3582 0.85 1.56 0.32±0.01 3629 0.46 0.21±0.06 7187 4000 1.80 1.40 0.30±0.14 4134 1.78 1.37 0.41±0.01 4111 1.63 0.41±0.06

Figures 3.2 and 3.3 provide a comparison of our re-derived optical Teff and log g with

APOGEE DR12, DR13 (SDSS Collaboration et al. 2016), and both ASPCAP and The

Cannon from DR14. A noticeable offset can be seen in Figure 3.2 with both DR14 sets for the same stars at ∼4000–5000 K, i.e. cool stars in this study tend to be cooler than

APOGEE. A similar offset can be seen in Figure 3.3 for the warmest and coolest stars in the sample. This is not entirely surprising given Fe II lines are unavailable in APOGEE

IR spectra, and indirect methods must be utilized for determining surface gravity.

31 5500

5000

4500 [K] (APOGEE) eff

T 4000 DR12 DR13 DR14 (The Cannon) DR14 (ASPCAP) 3500 3500 4000 4500 5000 5500 Teff [K] (This study)

Figure 3.2. Re-derived optical stellar temperatures compared to recent data releases from the APOGEE (DR12, gray; DR13, pink; DR14 (The Cannon), green; DR14 (ASPCAP), dark blue) IR survey for our calibration cluster stars. The dashed diagonal line represents perfect agreement. A global error bar is top left.

32 3.5

3.0

2.5 (APOGEE) g

log 2.0

DR12 1.5 DR13 DR14 (The Cannon) DR14 (ASPCAP) 1.5 2.0 2.5 3.0 3.5 log g (This study) Figure 3.3. Re-derived spectroscopic surface gravity compared to recent data releases from the APOGEE IR survey for our calibration cluster stars.. Plotting designations are the same as Figure 3.2.

33 3.2.3 Equivalent Width Analysis and Atomic Data

The goal of an equivalent width (EW) analysis is to relate the equivalent width (EWλ) to the column density number of atoms (Na) in a curve of growth regime. The relation is monotonic, but non-linear. There are three regimes depending on the optical depth at line center, τ0. When the optical depth is small, τ0 << 1, the regime is linear and

EWλ is proportional to Na. This is the linear part of the curve of growth, and returns an abundance which is considered reliable. When the optical depth is large, τ0 > 1, the regime is flat and returns EWλ proportional to the square root of the natural log of Na. Far from line center there is partial absorption and EWλ grows slowly with Na.

This is the flat part of the curve of growth. Unfortunately, many absorption features fall on this part of the curve of growth. If the optical depth is very large, the regime is damped, or “saturated”, and returns an abundance proportional to the square root of

Na. These are lines that are generally broad with deep flux, and which, unintuitively, return an abundance smaller than an abundance measured on the linear part of the curve of growth. Figure 3.4a illustrates the three regimes, while Figure 3.4b depicts the area of a spectral line measured below the continuum level, related to a rectangular line profile with the same area and equivalent width.

Continuum placement determines the area for measurement, and so, is a critical component to the EW measurement. The placement of the continuum can be global, i.e. across the entire spectrum, or local in the case of an obvious displacement caused by some vertical shifting, e.g. blending from a nearby line, noise and/or molecular contamination which can distort the continuum at a local level. In the case of line blanketing, or a

34 0.0

-0.5

∝ 1/2 -1.0 W Na )

λ ∝ 1/2 -1.5 W (ln Na)

W( ∝ W Na 10 -2.0 log -2.5

-3.0

-3.5 14 15 16 17 18 19 20 21 -2 log10 Na cm (b) Equivalent Width (a) Curve of Growth

Figure 3.4. Figure (a) illustrates the three curve of growth regimes; linear, flat and damped for measured EWs. (b) EWs are measured from the continuum and related to a rectangular line profile with the same area and equivalent width. decrease in the intensity of a stellar spectrum due to many closely spaced, unresolved lines, the continuum should be placed at the global level, or better, the line in question should be avoided entirely.

Abundances for fourteen species (Fe I & II, Na, Mg, Al, Si, Ca, Sc II, Ti I & II, V,

Ni, Zr I and Ba II) were derived via EW analyses. Lines were chosen for measurement both by consulting the Arcturus atlas (Hinkle et al. 2000) as well as a meticulous liter- ature search for previous studies of the same clusters, and in many instances the same stars. Given the high resolution of our spectra, I chose lines for analysis that were not expected to be severely blended with nearby features, or that suffered from line blanket- ing effects. Moreover, consistently strong or obviously blended lines were discarded by iterative processes described in the preceding section. Particular attention was given to lines stronger than 150 mA,˚ but in some cases EW measurements ∼160 mA˚ were allowed

35 due to the limited number of lines available for a particular species, or if they made little

change (. 0.02 dex) in the abundance outcomes. As an example, Ca lines tend to be strong, and can even be saturated in cool stars, so I have allowed a small number of EW measurements for these lines to exceed the 150 mA˚ ceiling for practical reasons, provided the abundances of these strong lines were self-consistent with abundances for other lines of the same species.

3.2.3.1 Equivalent Width Measurements

EW measurements were performed using two methods: First, a module in the automated

EW measuring program ROBOSPECT (Waters & Hollek 2013) was utilized for initial

EW measurements, and second, by using IRAF’s splot package for by-hand follow-up of

the initial measurements. A program was written to call ROBOSPECT, measure EWs

repetitively from an input line list, plot the output of best-fit measurements for each line,

and to create a sorted line list suitable for use with MOOG. For our flux normalized spec-

tra, the continuum placement was set to null i.e., continuum = 1, for measuring all lines.

This was a highly efficient method to determine which strong lines should be removed

from analysis, and to locate spectral lines directly affected by the infamous “picket fence”

described in §3.1. EW measurements by ROBOSPECT were accurate to less than ∼0.05 mA˚ for spectra with SNR & 100 when re-measured using the Gaussian fitting routine in

IRAF, but accuracy decreased by as much as ∼7–10 mA˚ with decreasing SNR of the spec- tra. Additionally, with a global continuum set to one, lines such as Mg λ6318 A,˚ which is affected by nearby Ca auto-ionization, require that a local continuum be determined (see

36 Culver 1967, for an example of the continuum displacement). Therefore, each line was

visually inspected for proper continuum placement, then carefully re-measured, making

use of the deblending function and Gaussian fitting routine in IRAF’s splot package.

3.2.3.2 Line List and Atomic Data

Accurate analyses of stellar abundances depend on quality atomic data suitable for the

stars in question. The line list was developed with the goal of deriving accurate abun-

dances given the Teff range for evolved stars in this study, and for the wavelength ranges

provided by observations. Where self-consistent I have used atomic data from recent

studies, otherwise adopted older, yet still reliable, oscillator strengths6. Atomic data for

Fe I were culled from recent lab studies by Ruffoni et al.(2014) and Den Hartog et al.

(2014), with the exception of λλ5775 & 6056 A˚ from O’Brian et al.(1991). For Fe II I

adopted oscillator strengths, referred to as gf-values, from Kurucz7 rather than the more

general oscillator strengths from the NIST database. The recently available atomic data

from Kurucz was also used to address hyperfine issues associated with Sc II, and utilized

for Si I with excitation potentials & 5.6 eV. For Si I with excitation potentials ∼5.1 eV

and lower, the “older” log gfs from the Kurucz database returned more self-consistent results. For Ni I, log gf values for λλ6176, 6204 and 6223 A˚ were taken from Wood et al.(2014); Fuhr et al.(1988) provided values for λλ5593 & 5760 A;˚ and for λ5805 A˚

6For an electronic transition to occur, an oscillating dipole must be induced by interaction of an atom’s or molecule’s electic field with electromagnetic radiation. In spectroscopy, an oscillator strength (f) refers to the probability that a transition will occur between energy levels in atoms or molecules. It is a dimensionless quantity often expressed on a logarithmic scale and with a statistical weight, g. 7See http://kurucz.harvard.edu/linelists.html, specifically the atomic data in gfnew (and gfall for the “older” data), and references contained therein.

37 Kurucz(2008) was adopted. Since Ba II is subject to hyperfine structure and/or isotopic

broadening, I utilized log gfs following McWilliam(1998). The final line lists with EW

measurements for all calibration cluster stars are presented in AppendixA.

3.2.3.3 Na, Sc and V

Na abundances can be affected by non-LTE (NLTE) effects, or a departure from statistical

equilibrium, and corrections are recommended for accurate abundance determinations.

NLTE effects for Na have a strong dependence on Teff , log g and lines chosen to derive

Na abundances (Mashonkina et al. 2000). Least affected are the λλ6154–6160 A˚ doublet

(see also Korotin & Komarov 1989), which I enlist for this study. Since corrections are expected to be small for our relatively cool, metal-rich stars, I apply no NLTE treatment to our Na abundances.

Odd-Z elements such as Scandium and Vanadium are subject to hyperfine structure

(hfs) stemming from nucleon-electron interactions, which can split absorption lines into multiple components and, in turn, impact chemical abundance analyses. In an effort to address these issues, I have incorporated log gfs for Sc II following recommendations by

Prochaska & McWilliam(2000) into our line list (see §3.2.3.2 for reference). Similarly, our input line list makes use of log gfs from a recent study by Lawler et al.(2014), specifically addressing hfs for V I lines selected for this study.

38 3.2.4 The Neutron Capture Elements

Most stable isotopes of elements heavier than the Fe-peak are produced through neutron

captures. Slow neutron capture, or the s-process, occurs when the timescale for neutron

capture is longer than the timescale for β decay. The reverse is true for the rapid, or

r-process, where the timescale for β decay is longer than the timescale for capture. The production site considered the “main” component for the majority of sprocess elements

(e.g., Ba and La) has been successfully modeled for low mass (M . 4M ) thermally

pulsing asymptotic giant branch (AGB) stars (see Busso et al. 1999; 2004, for reviews).

The primary neutron source in these low mass stars is provided by the 13C(α, n)16O

reaction arising from alternating ignition of H and He shells surrounding an inert C-O

core.

Conditions for r-element processes demand high neutron flux environments, and core

collapse supernovae (SNe) are plausible production site events (Hoyle 1954, Burbidge

et al. 1957). Yet theoretical SNe yield models face difficulties in reproducing the observ-

able r-process abundance distribution (see Sneden et al. 2008, and references therein).

Fortunately, recent observations of a blue “kilonova” (Metzger et al. 2010, Kasen et al.

2013, Tanaka & Hotokezaka 2013) from a nearby binary neutron star merger (GW170817)

offers a promising production site for r-process nuclei (Smartt et al. 2017). In short, the

ratio of the relative abundances of s-process to r-process elements trace the respective

contributions of AGB stellar winds and, quite possibly, Type II SNe and binary neu-

tron star mergers (see Metzger 2017, and references therein for a summary of the binary

neutron star merger event).

39 3.2.4.1 Zr and Ba

Zr I abundances have been derived via EW analysis from four lines in our spectra, λλ6127,

6134, 6140 and 6143 A,˚ although λ6140 A˚ is used only sparingly as it tends to be over-

abundant with respect to other Zr I lines. For Zr II I derive abundances from a single

line, λ6114 A,˚ by employing the spectrum synthesis driver synth for MOOG.

Three Ba II absorption lines are available for analysis in our wavelength range,

λλ5853, 6141 and 6496 A.˚ It is well known that Ba II λ6141 A˚ is terminally blended

with neutral Fe, and thus returns abundances that are unreliable (Hyland & Mould

1974, Mashonkina & Gehren 2000, Mishenina et al. 2013). Therefore, I have excluded

this line from our analysis, but include it in the line list for completeness. Ba II λ6496 A˚

is a deep and broad absorption line that is often saturated, and can also return unreliable

abundances. For this reason, I restrict our use of λ6496 A˚ to cases for which abundances

are unchanged by more than ± 0.05 dex from the remaining line, λ5853 A.˚

3.2.5 Spectrum Synthesis

While I employ EW measurements for Zr I and Ba II, all other n-capture element abun- dances were derived with the synth driver provided by MOOG (2017 version). Nearby molecular contamination were modeled and removed with detailed line lists kindly pro- vided by Chris Sneden (private communication). Three step-size sets of synthetic spectra were generated based on stellar parameters established via EW analysis, then modified to agree with the resolution of nearby absorption features. The code also allows for adjust-

40 ments in horizontal and vertical space, which were used in the case of small velocity or

continuum offsets of the observed spectra. The synthetic spectra were then best-fit to the

observed spectra by varying the input abundance. Uncertainty for the fit was determined

by resizing the window to show only the element of interest, essentially eliminating other

absorption lines from the fit. Figures 3.5 and 3.6 show examples of synthesis best-fit

for La II λ6390 A˚ and Eu II λ6645 A,˚ respectively, for M67 – 6425. The dotted line

represents the observed spectrum, with the best-fit synthetic spectrum shown as the blue

solid line. The step-size abundance, ±0.50, of the two remaining synthetically generated spectra are shown with corresponding red and green solid lines.

41 1.00

0.90

La +1.70 La +1.32 Relative Flux 0.80 La +0.70 M67 2M08514388+1156425 4655./ 2.50/ 0.00 vt = 1.50 6390.00 6390.25 6390.50 6390.75 6391.00 Wavelength (Å)

Figure 3.5. Spectrum synthesis for La II. The dotted line represents the observed spec- trum, with the best-fit synthetic spectrum shown as the blue solid line. The step-size abundance, ±0.50, and two remaining synthetically generated spectra are shown with corresponding red and green solid lines. Abundances are shown as log (A) values.

1.00

0.90

Eu +1.02 Eu +0.49 Relative Flux 0.80 Eu +0.02 M67 2M08514388+1156425 4655./ 2.50/ 0.00 vt = 1.50 6644.75 6645.00 6645.25 6645.50 Wavelength (Å)

Figure 3.6. Spectrum synthesis for Eu II. All plotting designations are the same as 3.5

The previous figures show synthesis for a higher SNR spectrum (SNR = 90) with

σ = ∼0.007 for the best-fits. Figure 3.7 shows a best-fit for Y II in a cooler star, with

a lower SNR spectrum (SNR = 75). Cooler stars are more likely to have molecular

contamination that affects the continuum and, combined with noise in the spectrum, can

42 distort the line making it asymmetric. For such cases, best-fits are determined both by-

eye and by lowest possible σ values. In this case, σ = 0.026. Synthetic best-fits σ > 0.04 have been discarded from analysis.

1.00

0.90

Y +2.95 Y +2.45 Relative Flux 0.80 Y +1.95 NGC 6791 2M19204971+3743426 3600./ 1.00/ +0.36 vt = 1.40 6794.75 6795.00 6795.25 6795.50 6795.75 6796.00 Wavelength (Å)

Figure 3.7. Spectrum synthesis for Y II. All plotting designations are the same as 3.5

3.2.5.1 Y II & Zr II

Y II and Zr II are representative of the first peak, light s-process (ls) elements. Yttrium

has only one stable isotope, 89Y, which has the “magic” number8 50 neutrons in its

nucleus. About 72% of Y is created by the s-process, with the r-process responsible for

the the remaining 28%9. For Y II I analyzed three absorption features, λλ5544, 5728

and 6795 A.˚ Considering there were two instrument setups that rendered available lines

on the either the bluer or redder ends of the spectrum, at least two Y II features were

available for analysis in every observed spectra.

8In nuclear physics, a magic number is a number of nucleons (either protons or neutrons, separately) arranged into complete shells within the atomic nucleus. Isotopes with a magic number of either protons or neutrons are especially stable. The seven most widely recognized magic numbers as of are 2, 8, 20, 28, 50, 82, and 126. 9All percentages of solar system abundances stated are based on Kappeler et al.(1989)

43 Zr has five stable isotopes, 90, 91, 92, 94 and 96, with 96 being the only isotope

made by r-process. The solar system s-fraction is ∼83%, leaving an r-fraction of ∼17%.

Only one reliable transition was employed for measuring Zr II abundances, λ6114 A.˚

Unfortunately this line fell inside the “picket fence” range for some stars, rendering it

inaccessible. Zr I and II abundances are in good agreement, with differences generally <

0.1 dex from star-to-star and cluster averages differing from ∼0.03–0.06 dex.

3.2.5.2 La II, Ce II & Nd II

La, Ce and Nd, along with Ba, belong to the second s-process peak, the heavy- s (hs)

elements. Lanthanum has only one stable isotope, 139La, with ∼75% of La created by

the s-process. Three transitions are utilized to derive La II abundances: λλ5805, 5880

and 6390 A,˚ which are available for analysis in all our spectra. All lines tend to be

unblended, but the continuum can be displaced by isotopes of CN. These were first

modeled by increasing the division factor to a large integer to assess changes in the

abundance and/or continuum. Once the isotope was known to be present, it could be

effectively removed from analysis. This technique was employed for every element derived

by spectrum synthesis.

There are four isotopes of Ce, but only two contribute significantly to solar system

abundances, 140Ce and 142Ce, ∼86% and 11%, respectively, of which ∼77% are produced

in s-process nucleosynthesis. Two Ce II lines were used for analysis, λλ5975 and 6043

A.˚ On occasion, like Zr II λ6114, the λ5975 line fell in the “picket fence” range of our

spectra and was summarily abandoned.

44 Neodymium (Z=60) has seven stable isotopes with 142Nd, an s-process only isotope,

contributing ∼27% to solar system abundances. Isotopes of Nd are almost evenly divided by r-process and s-process production. ∼48% of Nd isotopes are produced by the s- process and ∼52% by the r-process. Interestingly, Nd II tends to be slightly super-solar in Galactic evolved stars (e.g. see Pancino et al. 2010, Reddy et al. 2013a). There are a wealth of Nd transition features available in our optical data. Many, however tend to be blended with V I or other Fe-peak elements such as Cr I or Fe. For our analysis

I chose λλ5431, 5740, 5811 and 6740 A,˚ lines neither blended nor asymmetric. All Nd

II abundances were derived from three lines: λλ5431, 5740 and 5811 A˚ for our “bluer” spectra and λλ5740, 5811 and 6740 A˚ for the spectra from the red instrument setup.

3.2.5.3 Eu II

Eu is considered a “pure” r-process element. Both stable isotopes, 151Eu and 153Eu, which are present in nearly equal proportion, are almost entirely created by r-process events.

Only ∼2% of 153Eu is created by the s-process route, while ∼98% of the two long-lived isotopes are produced by the r-process. Eu abundances were determined from two lines in our spectra, λλ6437 and 6645 A.˚ While Eu λ6645 is widely used for Eu abundance analyses, and often the only abundance marker in optical abundance studies, I chose to include λ6437 in an effort to provide a more statistically viable Eu abundance. We adopted the line list from Lawler et al.(2001) for Eu II λλ6437 and 6645 to attenuate the effects of isotopic splitting, even while the study warns λ6437 is “severely blended with

Si I λ6437.71.” It is indeed asymmetric. To mitigate the effects of this blend, I chose a significantly lower log gf and accompanying excitation potential found in the Kurdatucz

45 database for the Si feature. To further mitigate the effect, I included the Si abundance derived from EW analysis for each star during spectrum synthesis. Using this method, the line-to-line variation in Eu abundance from λ6437 and λ6645 was < 0.05 dex.

3.3 Results and Discussion

Ten elements, including Fe, common to the APOGEE survey were analyzed by this study for abundances. Figure 3.8 shows the common elements, ∆[X/Fe] and ∆[Fe/H], as a function of ∆Teff so that perfect agreement would be the center of each plot. As the comparison to our re-derived temperatures showed (Figure 3.2), we can see the difference clearly in Figure 3.8. Our temperatures are somewhat cooler for stars compared to DR14 derived by ASPCAP and The Cannon, but are in good agreement with DR13. We also see some scatter in Na, Ti II and V. Other elements are with the common error, ± 0.1 dex.

46 0.4 DR13 DR14 (The Cannon) DR14 (ASPCAP) 0.2

0.0

[Fe/H] -0.2 ∆ -0.4

0.4 Na Mg Al

0.2

0.0 -0.2 -0.4 0.4

Si Ca Ti I 0.2

0.0

[X/Fe] -0.2 ∆ -0.4

0.4 Ti II V Ni 0.2

0.0 -0.2 -0.4 -100 0 100 -100 0 100 -100 0 100 ∆ Teff (K)

Figure 3.8. We show the common APOGEE elements, ∆[X/Fe] and ∆[Fe/H], as a func- tion of ∆Teff so that perfect agreement would be the center of each plot. A common error bar, ± 0.1 dex and ± 100 K, has been placed in the top left plot for reference.

47 3.4 Literature Comparisons

For multiple cluster studies, it is helpful to know how our results compare to other studies.

Each calibration cluster section includes a relevant literature compilation from previous

spectroscopic abundance studies for comparison to our mean cluster abundances. Errors

reported are the internal errors resulting from the line-to-line scatter and deviation from

the mean abundance within a measured species. At the end of each cluster section I also

include a table of star-by-star abundances derived for each cluster this study.

3.4.1 M67

M67 is one of the most well-studied open clusters in the Galactic disk. The references

to it in literature are too many to mention here. M67 is a solar age, ∼4.3 Gyr, nearby

open cluster (only ∼ 800 pcs) with notably solar metallicity. With its well populated

RGB, it’s bright red giants and He-core burning “red clump” stars have been the subject

of numerous studies (see Figure 3.1). M67 serves as a a basis for many multi-cluster

open studies because of its well known solar metallicity. This study is no different in that

respect. We have analyzed ten stars for this cluster and find our stellar parameters for all

M67 stars agree well with literature, and especially well with J11, for all stars common

to both studies. Temperatures are within ±60 K, and surface gravities are within ±0.20 cm s−1. We find excellent agreement with Pancino et al.(2010) for abundances, with the exception of Mg I (see Table 3.3). The study cites a difficulty with consistent atomic data for the discrepancy with Mg abundances, although it is consistent with a value found by Jacobson et al.(2011). For M67 I show an under-abundance for V I with respect to

48 other studies, but for other clusters I am in good agreement. It should be noted that

Al I abundances were derived from only six of the ten M67 stars due to a miscalculated grating setting for the SES in 2013.

For M67 I compare to eight other M67 abundance studies, presented in Table 3.3.

Neutron capture elements, from Y II–Eu II, are in very good agreement with literature values, especially the Pancino et al.(2010) study.

49 Table 3.3. Comparison of M67 mean cluster abundances with literature.

Ratio O17 a T00 b Y05 c P10 d F10 e J11 f R13 h O16 i

[FeI/H] −0.02 ± 0.06 −0.02 ± 0.06 0.02 ± 0.14 0.05 ± 0.02 0.03 ± 0.07 −0.01 ± 0.05 −0.08 ± 0.04 0.05 ± 0.04 [FeII/H] −0.01 ± 0.05 −0.02 ± 0.04 ...... −0.08 ± 0.05 . . . [Na/Fe] 0.06 ± 0.05 0.20 ± 0.12 0.30 ± 0.10 0.08 ± 0.09 0.15 ± 0.10 0.03 ± 0.06 0.25 ± 0.03 . . . [Mg/Fe] 0.04 ± 0.04 0.05 ± 0.14 0.16 ± 0.08 0.27 ± 0.04 0.09 ± 0.02 0.23 ± 0.07 0.16 ± 0.02 . . . [Al/Fe] 0.03 ± 0.05 0.09 ± 0.09 0.17 ± 0.05 0.03 ± 0.02 0.11 ± 0.07 . . . 0.09 ± 0.01 . . . [Si/Fe] 0.05 ± 0.07 0.07 ± 0.14 0.09 ± 0.11 0.10 ± 0.02 0.18 ± 0.04 0.21 ± 0.05 0.20 ± 0.02 . . . [Ca/Fe] −0.14 ± 0.07 0.00 ± 0.14 0.07 ± 0.06 −0.16 ± 0.03 −0.07 ± 0.02 −0.11 ± 0.07 0.04 ± 0.04 . . . [ScII/Fe] −0.03 ± 0.07 0.05 ± 0.09 . . . −0.03 ± 0.04 ...... 0.10 ± 0.04 . . . [Ti/Fe] −0.03 ± 0.07 0.03 ± 0.08 0.12 ± 0.07 −0.04 ± 0.06 −0.14 ± 0.05 −0.10 ± 0.05 −0.01 ± 0.04 . . . [TiII/Fe] −0.02 ± 0.05 ...... −0.13 ± ...... 0.01 ± 0.05 . . . [V/Fe] −0.03 ± 0.07 0.11 ± 0.08 . . . 0.15 ± 0.13 ...... 0.09 ± 0.05 . . . [Ni/Fe] −0.02 ± 0.06 0.01 ± 0.09 0.08 ± 0.10 0.05 ± 0.01 −0.02 ± 0.01 −0.01 ± 0.06 0.10 ± 0.03 . . . [YII/Fe] −0.04 ± 0.01 ...... −0.05 ± 0.04 ...... 0.03 ± 0.04 . . . [Zr/Fe] −0.10 ± 0.05 −0.15 ± 0.09 −0.28 ± 0.04 . . . −0.14 ± 0.09 −0.03 ± 0.06 −0.07 ± 0.05 . . . [ZrII/Fe] −0.11 ± 0.01 ...... −0.07 ± 0.03 . . . [BaII/Fe] 0.19 ± 0.01 −0.01 ± 0.02 −0.02 ± 0.05 0.25 ± 0.02 . . . 0.10 ± 0.07 g −0.16 ... [LaII/Fe] 0.11 ± 0.01 −0.10 ± . . . 0.11 ± 0.06 0.01 ± 0.03 . . . −0.15 ± 0.04 g 0.00 ± 0.03 . . . [CeII/Fe] 0.09 ± 0.01 0.09 ± 0.12 ...... −0.02 ... [NdII/Fe] 0.09 ± 0.01 ...... 0.08 ± 0.05 ...... 0.02 ± 0.04 0.00 ± 0.02 [EuII/Fe] 0.03 ± 0.01 0.08 ± . . . 0.06 ± 0.02 ...... −0.13 ± 0.07 g 0.08 −0.11 ± 0.04

a O17 = This study. Abundances derived via spectrum synthesis, and associated errors, shown in bold . The internal uncertainties for all literature comparisons are those quoted by the authors, where (± . . . ) represents abundances derived from a single line. b T00 = Tautvaiˇsieneet al.(2000,R ∼30,000 for Ce and Eu values, all others R ∼60,000) c Y05 = Yong et al.(2005) d P10 = Pancino et al.(2010, weighted averages) e F10 = Friel et al.(2010, lines less than 150 m A)˚ f J11 = Jacobson et al.(2011) g These values are taken from Jacobson & Friel(2013) h R13 = Reddy et al.(2013b, abundances in bold calculated by synthesis) i O16 = Overbeek et al.(2016). 50 For star-by-star abundance tables, I utilize the abbreviated APOGEE IDs noted in

Table 3.1. Table 3.4. Measured Abundances for M67.

Ratio 0076 1230 3520 4263 4264 5049

[FeI/H] −0.10±0.06 −0.01±0.05 −0.02±0.08 0.01±0.06 0.07±0.05 −0.02±0.07 [FeII/H] −0.10±0.04 −0.02±0.05 −0.02±0.06 0.00±0.05 0.06±0.07 0.00±0.06 [Na/Fe] 0.15±0.03 0.13±0.04 0.05±0.00 0.09±0.08 0.05±0.05 0.07±0.04 [Mg/Fe] 0.03±0.02 0.03±0.00 0.09±0.06 −0.01±0.04 0.01±0.02 0.03±0.07 [Al/Fe] 0.08±0.01 0.04±0.04 . . . 0.03±0.08 −0.02±0.06 . . . [Si/Fe] 0.12±0.09 0.01±0.08 0.09±0.06 −0.02±0.08 −0.11±0.05 0.07±0.07 [Ca/Fe] −0.11±0.08 −0.17±0.08 −0.19±0.09 −0.19±0.07 −0.10±0.04 −0.16±0.09 [ScII/Fe] 0.10±0.04 0.02±0.08 −0.02±0.03 −0.10±0.08 −0.09±0.09 −0.04±0.08 [Ti/Fe] 0.02±0.08 −0.06±0.04 −0.17±0.08 −0.07±0.07 0.00±0.09 0.01±0.09 [TiII/Fe] 0.04±0.06 −0.05±0.05 −0.15±0.06 −0.05±0.03 0.00±0.08 0.01±0.07 [V/Fe] 0.08±0.08 0.00±0.08 −0.12±0.05 0.00±0.08 0.08±0.05 −0.06±0.08 [Ni/Fe] 0.01±0.08 0.02±0.08 0.01±0.09 −0.03±0.04 −0.14±0.08 −0.02±0.05 [YII/Fe] 0.01±0.01 0.06±0.01 −0.04±0.01 −0.07±0.01 0.03±0.01 −0.02±0.01 [Zr/Fe] −0.10±0.00 −0.12±0.08 −0.10±0.10 −0.04±0.05 −0.09±0.03 −0.16±0.06 [ZrII/Fe] −0.01±0.01 −0.15±0.02 −0.18±0.00 −0.11±0.01 . . . −0.19±0.01 [BaII/Fe] 0.19±0.01 0.10±0.01 0.23±0.00 0.22±0.00 0.22±0.00 0.32±0.00 [LaII/Fe] 0.32±0.01 0.14±0.01 0.17±0.01 0.09±0.01 0.02±0.01 0.07±0.01 [CeII/Fe] 0.17±0.01 0.10±0.01 0.06±0.01 −0.03±0.00 0.09±0.01 0.01±0.01 [NdII/Fe] 0.28±0.01 0.07±0.01 0.07±0.01 0.12±0.01 0.09±0.01 0.05±0.01 [EuII/Fe] 0.06±0.01 0.04±0.01 0.12±0.01 −0.01±0.01 −0.04±0.01 0.01±0.01

51 3.4.2 NGC 188

NGC 188 is an old (∼7 Gyrs), sparsely populated open cluster in the direction of the

Pole star, Polaris. As it is circumpolar, or navigates around the North pole through the

night, it can be a difficult cluster to observe. For this cluster, I observed only three stars.

For clusters included in this study I show a solar scaled abundance, or slight under-

abundance, for Zr I. For NGC 188 this compares poorly with Friel et al.(2010), for which

I have no stars in common. However, two stars common to this study and J11 for NGC

188, Zr I abundances are in good agreement. For J11, N188 – 7215, [Zr/Fe] = −0.13 (σ

= 0.26 dex), which is N188 – 7563, [Zr/Fe] = −0.18 ± 0.02 dex, for this study. For J11

N188 – 6712, [Zr/Fe] = −0.31 (σ = 0.35 dex), comparable to our N188 – 6329, [Zr/Fe]

= −0.20 ± 0.05 dex.

However, for Friel et al.(2010), the spread in Zr I abundances is large, 0.05 dex for four stars, compared to this study, 0.04 dex for three stars. For 28 stars in J11 the spread is far greater at 0.93 dex. Figure 3.9 shows two synthetic spectra generated from the adopted stellar parameters for NGC 188 – 3336 and fit to the observed spectrum, shown as a dotted black line. The red synthetic spectrum was created to show a solar zirconium metallicity, [Zr/Fe] = 0, while the gray synthetic spectrum was generated to show our derived abundance, [Zr/Fe] = −0.18 dex. For each Zr I line used in this analysis, marked at λλ6127, 6134 ad 6143, the red synthetic spectrum shows an over-abundance at solar metallicity. A good fit was found with our derived Zr I abundance for this star,

−0.18 dex, which reflects our mean Zr I cluster abundance for NGC 188 (see Table 3.5).

52 1.00

0.80

0.60 Zr Zr Zr 0.40 [Zr/Fe] = 0.00 Relative Flux [Zr/Fe] = −0.18 0.20 NGC 188 2M00463920+8523336 4300./ 2.00/ 0.06 vt = 1.60 6126 6128 6130 6132 6134 6136 6138 6140 6142 6144 Wavelength (Å)

Figure 3.9. Synthetic spectra generated for NGC 188 – 3336 in the wavelength range for all Zr I lines used in this study. The observed spectrum is shown as a dotted black line, with synthetic spectra as gray and red solid lines. The gray synthetic spectrum represents the best fit for the star’s stellar parameters, labeled at the bottom right, and our derived abundance for this star, [Zr/Fe] = −0.18 dex.

Table 3.5. Comparison of N188 mean cluster abundances with literature.

Ratio O17 a F10 b J11 c O16 e

[FeI/H] 0.04 ± 0.07 0.12 ± 0.02 −0.03 ± 0.04 0.12 ± 0.04 [FeII/H] 0.03 ± 0.05 ...... [Na/Fe] 0.24 ± 0.05 0.15 ± 0.03 0.10 ± 0.05 . . . [Mg/Fe] 0.00 ± 0.07 0.17 ± 0.10 0.26 ± 0.05 . . . [Al/Fe] 0.08 ± 0.07 0.20 ± 0.03 ...... [Si/Fe] −0.05 ± 0.07 0.17 ± 0.08 0.25 ± 0.05 . . . [Ca/Fe] −0.07 ± 0.06 −0.04 ± 0.08 −0.04 ± 0.06 . . . [ScII/Fe] −0.21 ± 0.07 ...... [Ti/Fe] 0.02 ± 0.06 0.05 ± 0.12 0.14 ± 0.05 . . . [TiII/Fe] 0.05 ± 0.04 ...... [V/Fe] 0.00 ± 0.07 ...... [Ni/Fe] −0.02 ± 0.07 −0.03 ± 0.05 0.08 ± 0.05 . . . [YII/Fe] −0.02 ± 0.02 ...... Continued on Next Page. . .

53 Table 3.5 – Continued

Ratio O17 a F10 b J11 c O16 e

[Zr/Fe] −0.18 ± 0.03 0.15 ± 0.23 0.10 ± 0.05 . . . [ZrII/Fe] −0.21 ± 0.01 ...... [BaII/Fe] −0.05 ± 0.02 . . . 0.04 ± 0.17 d ... [LaII/Fe] −0.05 ± 0.02 ... −0.06 ± 0.22 d ... [CeII/Fe] 0.09 ± 0.01 ...... [NdII/Fe] 0.07 ± 0.01 ...... 0.14 ± 0.05 [EuII/Fe] −0.01 ± 0.01 ... −0.18 ± 0.12 d −0.12 ± 0.07 a O17 = This study. Abundances derived via spectrum synthesis, and associated errors, shown in bold. b F10 = Friel et al.(2010, lines less than 150m A)˚ c J11 = Jacobson et al.(2011) d These values are taken from Jacobson & Friel(2013) e O16 = Overbeek et al.(2016).

Table 3.6. Measured Abundances for NGC 188.

Ratio 3336 6329 7563

[FeI/H] 0.06±0.07 0.02±0.07 0.04±0.07 [FeII/H] 0.06±0.06 0.00±0.04 0.04±0.04 [Na/Fe] 0.31±0.06 0.14±0.06 0.26±0.03 [Mg/Fe] −0.07±0.09 0.02±0.03 0.05±0.08 [Al/Fe] 0.10±0.09 0.06±0.05 0.08±0.08 [Si/Fe] −0.04±0.08 −0.05±0.07 −0.06±0.07 [Ca/Fe] −0.07±0.08 −0.12±0.08 −0.03±0.03 [ScII/Fe] −0.22±0.08 −0.23±0.07 −0.18±0.08 [Ti/Fe] 0.03±0.08 −0.02±0.05 0.04±0.06 [TiII/Fe] 0.08±0.04 −0.02±0.05 0.09±0.02 [V/Fe] −0.02±0.04 0.00±0.08 0.03±0.08 [Ni/Fe] −0.03±0.07 −0.03±0.07 −0.01±0.06 [YII/Fe] −0.03±0.01 −0.04±0.02 0.01±0.02 [Zr/Fe] −0.16±0.03 −0.20±0.05 −0.18±0.02 [ZrII/Fe] . . . −0.21±0.00 . . . [BaII/Fe] −0.04±0.00 −0.04±0.00 −0.07±0.05 [LaII/Fe] −0.06±0.01 −0.03±0.02 −0.05±0.02 [CeII/Fe] −0.04±0.01 0.11±0.01 0.21±0.02 [NdII/Fe] 0.06±0.01 0.07±0.01 0.08±0.02 [EuII/Fe] −0.05±0.02 −0.02±0.01 0.03±0.01

54 3.4.3 NGC 7789

Another open cluster with a well-populated red giant branch is NGC 7789. NGC 7789 is a nearby open cluster at ∼2.3 kpc making it a good choice for abundance studies. For this study I have selected five giants along the red giant branch.

Again, I note that Al I has been neglected for this cluster due to the unfortunate grating setting mentioned earlier. For NGC 7789, as for M67, I find excellent agreement with our abundances of n-capture elements with the Pancino et al.(2010) study.

Table 3.7. Comparison of NGC 7789 mean cluster abundances with literature.

Ratio O17 a T05 b P10 c J11 d O16 e

[FeI/H] −0.03 ± 0.06 −0.04 ± 0.05 0.05 ± 0.01 0.02 ± 0.04 0.00 ± 0.03 [FeII/H] −0.02 ± 0.05 −0.04 ± 0.05 ...... [Na/Fe] 0.11 ± 0.06 0.28 ± 0.07 −0.05 ± 0.13 0.09 ± 0.05 . . . [Mg/Fe] −0.07 ± 0.04 0.18 ± 0.07 0.22 ± 0.07 0.14 ± 0.05 . . . [Al/Fe] . . . 0.18 ± 0.08 −0.03 ± 0.09 ...... [Si/Fe] 0.08 ± 0.06 0.14 ± 0.05 −0.01 ± 0.02 0.25 ± 0.05 . . . [Ca/Fe] −0.17 ± 0.08 0.14 ± 0.07 −0.18 ± 0.09 0.01 ± 0.05 . . . [ScII/Fe] −0.03 ± 0.06 −0.02 ± 0.07 0.08 ± 0.02 ...... [Ti/Fe] −0.02 ± 0.06 −0.03 ± 0.07 −0.03 ± 0.09 −0.05 ± 0.04 . . . [TiII/Fe] −0.01 ± 0.05 ...... [V/Fe] 0.03 ± 0.08 0.09 ± 0.12 −0.01 ± 0.09 ...... [Ni/Fe] 0.01 ± 0.06 −0.02 ± 0.05 −0.01 ± 0.01 0.00 ± 0.05 . . . [YII/Fe] 0.01 ± 0.01 0.13 ± 0.13 0.08 ± 0.09 ...... [Zr/Fe] −0.08 ± 0.06 −0.02 ± 0.13 . . . 0.18 ± 0.05 . . . [ZrII/Fe] −0.11 ± 0.01 ...... [BaII/Fe] 0.40 ± 0.00 . . . 0.45 ± 0.05 ...... [LaII/Fe] 0.13 ± 0.01 . . . 0.11 ± 0.05 ...... [CeII/Fe] 0.00 ± 0.01 0.09 ± 0.13 ...... [NdII/Fe] 0.14 ± 0.01 . . . 0.17 ± 0.30 . . . 0.16 ± 0.03 [EuII/Fe] 0.02 ± 0.01 0.02 ± 0.12 ...... −0.11 ± 0.03

a O17 = This study. Abundances derived via spectrum synthesis, and associated errors, shown in bold. b T05 = Tautvaiˇsien˙eet al.(2005). c P10 = Pancino et al.(2010, weighted averages). d J11 = Jacobson et al.(2011). e O16 = Overbeek et al.(2016).

55 Table 3.8. Measured Abundances for NGC 7789.

Ratio 0125 0586 1221 1459 7208

[FeI/H] −0.07±0.06 −0.06±0.05 0.02±0.06 −0.02±0.06 −0.03±0.05 [FeII/H] −0.06±0.07 −0.04±0.04 0.02±0.06 −0.02±0.06 0.00±0.04 [Na/Fe] 0.12±0.04 0.21±0.05 0.03±0.07 0.08±0.06 0.09±0.07 [Mg/Fe] −0.08±0.03 0.14±0.00 −0.20±0.02 −0.11±0.08 −0.12±0.08 [Al/Fe] ...... [Si/Fe] 0.11±0.06 0.13±0.05 −0.01±0.05 0.13±0.07 0.05±0.06 [Ca/Fe] −0.24±0.08 −0.15±0.09 −0.24±0.07 −0.11±0.10 −0.10±0.09 [ScII/Fe] −0.04±0.01 −0.07±0.08 −0.03±0.06 0.14±0.06 −0.13±0.07 [Ti/Fe] −0.02±0.07 0.01±0.05 0.00±0.08 −0.02±0.06 −0.08±0.06 [TiII/Fe] −0.02±0.05 0.01±0.08 0.01±0.04 0.00±0.03 −0.07±0.03 [V/Fe] 0.05±0.08 0.05±0.06 0.01±0.09 0.03±0.09 0.00±0.09 [Ni/Fe] 0.00±0.06 0.02±0.05 0.02±0.05 0.00±0.06 −0.01±0.07 [YII/Fe] −0.07±0.02 −0.05±0.02 0.09±0.01 −0.03±0.01 0.03±0.01 [Zr/Fe] −0.10±0.06 −0.03±0.07 −0.09±0.04 −0.14±0.05 −0.03±0.06 [ZrII/Fe] −0.17±0.02 . . . 0.00±0.01 . . . −0.16±0.01 [BaII/Fe] 0.52±0.00 0.28±0.00 0.30±0.00 0.41±0.00 0.47±0.00 [LaII/Fe] 0.13±0.01 0.15±0.02 0.18±0.02 0.09±0.01 0.08±0.01 [CeII/Fe] −0.11±0.02 0.04±0.01 −0.12±0.01 0.08±0.01 0.09±0.01 [NdII/Fe] 0.17±0.01 0.09±0.02 0.11±0.01 0.17±0.01 0.16±0.01 [EuII/Fe] 0.06±0.01 −0.04±0.01 0.14±0.01 −0.03±0.01 −0.03±0.01

3.4.4 NGC 2420

NGC 2420 is an older, more distant open cluster residing outside the solar neighborhood.

At [Fe/H]∼−0.20 dex, NGC 2420 is set apart from other open clusters primarily for its metal-deficiency with respect to other nearby clusters. For this cluster I chose six

RGB stars for abundance analysis and confirm it is a slightly metal-poor open cluster, consistent with other studies.

56 Table 3.9. Comparison of NGC 2420 mean cluster abundances with literature.

Ratio O17 a P10 b J11 c

[FeI/H] −0.18 ± 0.06 −0.05 ± 0.03 −0.20 ± 0.06 [FeII/H] −0.18 ± 0.05 ...... [Na/Fe] 0.10 ± 0.03 −0.04 ± 0.07 0.08 ± 0.06 [Mg/Fe] −0.02 ± 0.07 0.09 ± 0.06 0.11 ± 0.09 [Al/Fe] 0.01 ± 0.06 −0.10 ± 0.01 . . . [Si/Fe] 0.03 ± 0.06 0.04 ± 0.01 0.21 ± 0.07 [Ca/Fe] −0.03 ± 0.05 −0.09 ± 0.02 0.10 ± 0.07 [ScII/Fe] −0.02 ± 0.08 0.07 ± 0.05 . . . [Ti/Fe] 0.13 ± 0.07 −0.03 ± 0.02 −0.08 ± 0.08 [TiII/Fe] 0.14 ± 0.04 ...... [V/Fe] −0.01 ± 0.07 −0.05 ± 0.13 . . . [Ni/Fe] −0.02 ± 0.06 0.00 ± 0.02 −0.01 ± 0.07 [YII/Fe] −0.06 ± 0.06 0.05 ± 0.08 . . . [Zr/Fe] 0.02 ± 0.05 . . . 0.24 ± 0.08 [ZrII/Fe] −0.04 ± 0.06 ...... [BaII/Fe] 0.37 ± 0.06 0.57 ± 0.02 . . . [LaII/Fe] 0.11 ± 0.06 0.23 ± 0.09 . . . [CeII/Fe] 0.14 ± 0.06 ...... [NdII/Fe] 0.20 ± 0.06 0.19 ± 0.17 . . . [EuII/Fe] 0.04 ± 0.06 ......

a O17 = This study. Abundances derived via spectrum synthe- sis, and associated errors, shown in bold. b P10 = Pancino et al.(2010, weighted averages) c J11 = Jacobson et al.(2011).

Table 3.10. Measured Abundances for NGC 2420.

Ratio 1418 4589 5050 5508 8015 8244

[FeI/H] −0.18±0.07 −0.29±0.05 −0.15±0.07 −0.09±0.06 −0.23±0.05 −0.16±0.05 [FeII/H] −0.17±0.03 −0.29±0.07 −0.15±0.07 −0.09±0.03 −0.23±0.06 −0.15±0.06 [Na/Fe] 0.18±0.03 0.16±0.05 0.03±0.01 0.08±0.07 0.07±0.01 0.05±0.02 [Mg/Fe] −0.04±0.09 0.00±0.06 0.04±0.08 0.03±0.06 −0.04±0.08 −0.08±0.05 [Al/Fe] −0.02±0.09 0.08±0.09 −0.03±0.01 −0.01±0.05 0.02±0.03 0.01±0.09 [Si/Fe] 0.05±0.08 0.05±0.07 0.05±0.07 −0.09±0.08 0.10±0.05 0.00±0.04 [Ca/Fe] 0.07±0.05 −0.09±0.03 −0.01±0.05 −0.03±0.06 0.03±0.02 −0.12±0.07 [ScII/Fe] 0.13±0.07 −0.20±0.09 0.05±0.06 −0.20±0.09 0.07±0.09 0.02±0.07 [Ti/Fe] 0.12±0.06 0.23±0.04 0.10±0.07 0.11±0.09 0.21±0.07 0.02±0.06 [TiII/Fe] 0.12±0.02 0.25±0.06 0.10±0.03 0.12±0.05 0.20±0.05 0.04±0.03 [V/Fe] −0.05±0.06 0.13±0.07 −0.01±0.05 −0.09±0.08 −0.01±0.08 −0.03±0.07 [Ni/Fe] −0.06±0.09 −0.03±0.06 −0.01±0.08 −0.01±0.03 0.00±0.08 −0.03±0.04 [YII/Fe] −0.18±0.02 −0.04±0.01 −0.06±0.01 0.02±0.03 0.00±0.01 −0.08±0.02 [Zr/Fe] 0.05±0.07 −0.04±0.04 0.05±0.06 0.16±0.05 −0.10±0.04 0.02±0.04 Continued on Next Page. . .

57 Table 3.10 – Continued

Ratio 1418 4589 5050 5508 8015 8244

[ZrII/Fe] −0.01±0.01 −0.22±0.00 0.08±0.01 0.03±0.03 −0.12±0.01 0.00±0.02 [BaII/Fe] 0.40±0.00 0.39±0.00 0.28±0.08 0.45±0.00 0.47±0.00 0.23±0.01 [LaII/Fe] 0.15±0.01 0.10±0.01 0.10±0.02 0.05±0.02 0.14±0.02 0.10±0.01 [CeII/Fe] 0.10±0.02 0.22±0.01 0.12±0.02 0.07±0.01 0.13±0.02 0.17±0.02 [NdII/Fe] 0.34±0.01 0.30±0.01 0.16±0.02 0.23±0.03 0.09±0.01 0.09±0.01 [EuII/Fe] 0.13±0.01 0.04±0.01 0.03±0.01 0.04±0.02 0.05±0.01 −0.03±0.01

3.4.5 NGC 6819

For NGC 6819, I show an almost solar abundance for all elements considered in this study with the exception of Na, for which Bragaglia et al.(2001) also found an over-abundance.

The greatest discrepancy between the two studies is with Ba II (see Table 3.11). This discrepancy might be explained with the differences in solar values and lines chosen to measure the Ba abundance for each study. The Bragalia study employs only one line for measuring Ba abundances, λ6496, which can often have broad wings and deep flux, and consequently can often return an unreliable abundance determination. They also utilize a much larger solar Ba abundance than this study, log(Ba) = 2.34 as opposed to 2.18 for this study. Further, for three stars in their study, only one abundance determination for Ba is reported. Our individual star abundances can be found in Table 3.12.

58 Table 3.11. Comparison of NGC 6819 mean cluster abundances with literature.

Ratio O17 a B01 b C17 c

[FeI/H] 0.01 ± 0.07 0.09 ± 0.03 . . . [FeII/H] −0.01 ± 0.07 0.09 ± 0.03 . . . [Na/Fe] 0.26 ± 0.02 0.47 ± 0.07 . . . [Mg/Fe] 0.01 ± 0.04 −0.12 ± 0.07 . . . [Al/Fe] 0.04 ± 0.02 −0.07 ± 0.07 . . . [Si/Fe] 0.03 ± 0.07 0.18 ± 0.04 . . . [Ca/Fe] −0.05 ± 0.08 −0.04 ± 0.06 . . . [ScII/Fe] −0.08 ± 0.05 0.04 ± 0.04 . . . [Ti/Fe] 0.05 ± 0.07 −0.01 ± 0.04 . . . [TiII/Fe] 0.05 ± 0.03 ...... [V/Fe] 0.05 ± 0.06 −0.09 ± 0.07 . . . [Ni/Fe] −0.01 ± 0.07 0.01 ± 0.02 . . . [YII/Fe] −0.02 ± 0.01 0.16 ± 0.10 . . . [Zr/Fe] −0.05 ± 0.05 ...... [ZrII/Fe] 0.00 ± 0.01 ...... [BaII/Fe] 0.11 ± 0.00 −0.27 ± 0.15 . . . [LaII/Fe] 0.14 ± 0.02 ...... [CeII/Fe] 0.01 ± 0.02 ... −0.08 ± 0.06 [NdII/Fe] 0.16 ± 0.01 ...... [EuII/Fe] 0.02 ± 0.01 ......

a O17 = This study. Abundances derived via spectrum synthe- sis, and associated errors, shown in bold. b B01 = Bragaglia et al.(2001). c C17 = Cunha et al.(2017, from APOGEE IR spectra).

Table 3.12. Measured Abundances for NGC 6819.

Ratio 0532 2111 2283

[Fe/H] −0.04±0.07 −0.02±0.07 0.08±0.07 [FeII/H] −0.05±0.08 −0.02±0.08 0.05±0.06 [Na/Fe] 0.19±0.01 0.33±0.04 0.25±0.02 [Mg/Fe] 0.16±0.04 0.00±0.06 −0.13±0.01 [Al/Fe] 0.23±0.01 0.04±0.06 −0.14±0.00 [Si/Fe] 0.00±0.05 0.02±0.08 0.06±0.08 [Ca/Fe] 0.01±0.07 −0.02±0.08 −0.13±0.08 [ScII/Fe] 0.00±0.06 −0.12±0.09 −0.12±0.01 [Ti/Fe] 0.05±0.08 0.05±0.08 0.06±0.07 [TiII/Fe] 0.05±0.02 0.05±0.01 0.06±0.06 [V/Fe] 0.07±0.07 0.06±0.09 0.03±0.03 [Ni/Fe] 0.00±0.06 0.00±0.08 −0.04±0.05 [YII/Fe] −0.11±0.01 0.09±0.02 −0.04±0.01 [Zr/Fe] 0.12±0.06 −0.02±0.06 −0.24±0.03 [ZrII/Fe] 0.11±0.01 −0.10±0.02 . . . Continued on Next Page. . .

59 Table 3.12 – Continued

Ratio 0532 2111 2283

[BaII/Fe] 0.06±0.01 0.14±0.00 0.14±0.00 [LaII/Fe] 0.13±0.02 0.16±0.02 0.12±0.01 [CeII/Fe] −0.05±0.02 0.06±0.02 0.03±0.01 [NdII/Fe] 0.17±0.00 0.11±0.00 0.20±0.01 [EuII/Fe] −0.02±0.02 0.20±0.02 −0.12±0.01

3.4.6 NGC 6791

NGC 6791 is an old, distant open cluster that is also oddly metal-rich for the Galactic plane. Our measured abundances for NGC 6791 are consistent with those found in literature, including a large spread from study to study with barium.

Table 3.13. Comparison of NGC 6791 mean cluster abundances with literature.

Ratio O17 a C06 b O06 c C07 d

[FeI/H] 0.33 ± 0.09 0.37 ± 0.06 0.35 ± 0.08 0.47 ± 0.07 [FeII/H] 0.34 ± 0.06 ...... [Na/Fe] 0.10 ± 0.06 ...... 0.13 ± 0.21 [Mg/Fe] −0.07 ± 0.04 . . . −0.03 ± 0.09 0.20 ± 0.05 [Al/Fe] −0.11 ± 0.04 −0.15 ± . . . 0.05 ± 0.15 −0.21 ± 0.09 [Si/Fe] 0.04 ± 0.05 0.02 ± 0.03 0.02 ± 0.19 −0.01 ± 0.10 [Ca/Fe] −0.05 ± ... −0.03 ± 0.08 −0.03 ± 0.09 −0.15 ± 0.08 [ScII/Fe] −0.15 ± 0.05 ...... −0.13 ± 0.07 [Ti/Fe] −0.06 ± 0.06 −0.02 ± 0.04 0.03 ± 0.15 0.03 ± 0.09 [TiII/Fe] −0.07 ± 0.04 ...... [V/Fe] −0.06 ± 0.09 ...... [Ni/Fe] −0.04 ± 0.04 0.00 ± 0.14 . . . −0.07 ± 0.07 [YII/Fe] −0.09 ± 0.02 ...... [Zr/Fe] −0.15 ± 0.07 ...... [ZrII/Fe] −0.21 ± 0.01 ...... [BaII/Fe] −0.07 ± ... −0.13 ± ...... 0.28 ± 0.12 [LaII/Fe] 0.10 ± 0.02 ...... [CeII/Fe] 0.12 ± 0.01 ...... [NdII/Fe] 0.02 ± 0.01 ...... [EuII/Fe] −0.13 ± 0.02 ...... −0.17 ± 0.00

a O17 = This study. Abundances derived via spectrum synthesis, and associated errors, shown in bold. b C06 = Carraro et al.(2006) c O06 = Origlia et al.(2006) d C07 = Carretta et al.(2007).

60 Table 3.14. Measured Abundances for NGC 6791.

Ratio 7187 3426

[FeI/H] 0.30±0.14 0.36±0.04 [FeII/H] 0.31±0.06 0.37±0.05 [Na/Fe] 0.13±0.02 0.07±0.08 [Mg/Fe] −0.08±0.03 −0.06±0.04 [Al/Fe] −0.10±0.02 −0.12±0.05 [Si/Fe] 0.07±0.03 0.01±0.06 [Ca/Fe] −0.05±0.00 −0.04±0.00 [ScII/Fe] −0.17±0.03 −0.13±0.06 [Ti/Fe] −0.08±0.04 −0.03±0.07 [TiII/Fe] −0.08±0.04 −0.05±0.00 [V/Fe] 0.02±0.11 −0.14±0.07 [Ni/Fe] −0.03±0.04 −0.04±0.03 [YII/Fe] −0.07±0.01 −0.12±0.02 [Zr/Fe] −0.10±0.08 −0.19±0.05 [ZrII/Fe] −0.23±0.01 −0.18±0.01 [BaII/Fe] −0.03±0.00 −0.10±0.00 [LaII/Fe] 0.13±0.02 0.07±0.02 [CeII/Fe] 0.14±0.00 0.09±0.01 [NdII/Fe] −0.03±0.01 0.06±0.01 [EuII/Fe] −0.07±0.02 −0.18±0.02

We summarize the above sections in Figures 3.10 and 3.11 with a cluster-to-cluster comparison to literature. In each plot the dotted line indicates perfect agreement with this study.

In Figure 3.11, M67 is represented by dark blue filled circles, NGC 188 by light blue.

NGC 2420 is depicted by the black filled circles, while NGC 6791, NGC 6819 and NGC

7789 are shown as orange, red and pink filed circles, respectively. For more specific values, consult the literature compilation tables above.

61 Figure 3.10. ∆[Fe/H] cluster-to-cluster comparison to literature. The dotted line indi- cates perfect agreement with this study.

62 Na Mg Al

Si Ca Sc II

Ti I Ti II V

Ni Y II Zr I

Zr II Ba II La II

Ce II Nd II Eu II

Figure 3.11. [X/Fe] cluster-to-cluster comparison to literature. M67 is represented by dark blue filled circles, NGC 188 by light blue. NGC 2420 is depicted by the black filled circles, while NGC 6791, NGC 6819 and NGC 7789 are shown as orange, red and pink filed circles, respectively.

63 3.4.7 Sources of Uncertainty

We estimate the uncertainties in stellar parameters to be ± 100 K in Teff , 0.3 dex in log g, and 0.3 km s−1 in microturbulence. Due to the range of stellar parameters in this study, I have determined uncertainties in abundances, relative to hydrogen, for the warmest and coolest stars in this study by varying the parameters by the given uncertainties.

We provide Table 3.15 as a guide to uncertainties due to variations in stellar param- eters for the warmest and coolest stars in this study with the highest SNR.

Table 3.15. Abundance sensitivities to our re-derived model atmospheres for the warmest and coolest stars in our sample with the highest SNR. Entries show the effects on abun- dances relative to hydrogen, ∆[A/H].

Element ∆Teff ± 100 ∆log g ± 0.30 ∆ξ ± 0.30

(K) (cgs) (dex)

M67 – 9493, ∼5160 K

Fe I ±0.06 ±0.00 ±0.09

Fe II ±0.09 ±0.12 ±0.08

Na I ±0.06 ±0.02 ±0.04

Mg I ±0.04 ±0.04 ±0.07

Al I ±0.05 ±0.02 ±0.04

Si I ±0.02 ±0.03 ±0.03

Ca I ±0.09 ±0.04 ±0.12

Sc II ±0.02 ±0.11 ±0.07

Ti I ±0.11 ±0.00 ±0.04

Ti II ±0.03 ±0.11 ±0.04

VI ±0.13 ±0.00 ±0.05

Continued on Next Page. . .

64 Table 3.15 – Continued

Element ∆Teff ± 100 ∆log g ± 0.30 ∆ξ ± 0.30

(K) (cgs) (dex)

Ni I ±0.02 ±0.03 ±0.07

YII ±0.00 ±0.03 ±0.01

Zr I ±0.15 ±0.01 ±0.01

Zr II ±0.01 ±0.08 ±0.02

Ba II ±0.02 ±0.07 ±0.27

La II ±0.00 ±0.06 ±0.00

Ce II ±0.00 ±0.04 ±0.00

Nd II ±0.02 ±0.09 ±0.02

Eu II ±0.00 ±0.08 ±0.01

NGC 6791 – 7187, ∼4000 K

Fe I ±0.05 ±0.07 ±0.18

Fe II ±0.20 ±0.19 ±0.06

Na I ±0.06 ±0.02 ±0.16

Mg I ±0.05 ±0.02 ±0.12

Al I ±0.04 ±0.01 ±0.12

Si I ±0.12 ±0.08 ±0.07

Ca I ±0.07 ±0.03 ±0.23

Sc II ±0.04 ±0.13 ±0.18

Ti I ±0.08 ±0.04 ±0.20

Ti II ±0.06 ±0.14 ±0.06

VI ±0.09 ±0.06 ±0.31

Ni I ±0.07 ±0.08 ±0.13

YII ±0.00 ±0.04 ±0.07

Continued on Next Page. . .

65 Table 3.15 – Continued

Element ∆Teff ± 100 ∆log g ± 0.30 ∆ξ ± 0.30

(K) (cgs) (dex)

Zr I ±0.12 ±0.06 ±0.28

Zr II ±0.02 ±0.10 ±0.09

Ba II ±0.00 ±0.10 ±0.36

La II ±0.01 ±0.05 ±0.09

Ce II ±0.00 ±0.04 ±0.11

Nd II ±0.02 ±0.07 ±0.08

Eu II ±0.01 ±0.04 ±0.10

3.5 SUMMARY

We have determined abundances for twelve elements via EW measurements (Fe I & II,

Na, Mg, Al, Si, Ca, Sc II, Ti I & II, V, Ni, Zr I and Ba II) and six by spectral synthesis

(Y II, Zr II, La II, Ce II, Nd II and Eu II) for six well studied open clusters. We find our abundance determinations for all six open clusters to be reasonably consistent with values found in literature. Interestingly, in Figure 3.12 we again see the large spread in Ba abundances for the six clusters under consideration. Our measurements for Ba abundances agree reasonably well with those found in literature, so we can assume that the large spread for Ba II in open clusters is real and must have a physical origin.

66 1.0

0.8

0.6

0.4

0.2 [X/Fe] 0.0

-0.2

-0.4

-0.6 [Fe/H] Na Mg Al Si Ca Sc Ti V Ni Y Zr Ba La Ce Nd Eu

Figure 3.12. Boxplot of elements derived in this study for six clusters. The designations for plotting are the same as Figure 2.1

This study shows Zr to be slightly under-abundant compared to values found in litera- ture, which can also be seen in Figure 3.12. With the rather large range of temperatures for stars in this study, I find a slight temperature dependence for both Zr I & Zr II abundances.

67 0.4

0.20 0.3

0.10 0.2

0.00 0.1 [Fe/H] [Zr I,II/Fe] -0.10 0.0

-0.20 -0.1

-0.2 3600 3800 4000 4200 4400 4600 4800 5000 5200 Teff (K)

Figure 3.13. Star-by-star plot for mean Zr I and Zr II abundances as a function of tem- perature. The dashed line represents solar values. Super solar abundances for zirconium appear for only the warmest stars in this sample.

68 However, the same temperature dependence does not appear for Ba II. Ba II as function of temperature appears to be fairly flat, with little dependence on metallicity.

0.4 0.20 0.10 0.00 0.3 -0.10 [Zr I,II/Fe] -0.20 0.2 3800 4200 4600 5000 0.1 [Fe/H] 0.60 0.40 0.0 0.20

[Ba/Fe] 0.00 -0.1 -0.20 3800 4200 4600 5000 -0.2 Teff (K)

Figure 3.14. Zr and Ba plotted as function of temperature, with a dashed line representing solar values. No correlation with temperature can found for Ba with these stars, where Zr appears to be super solar for only the warmest stars.

69 For stars in this limited sample, Zr shows no correlation to metallicity. It is also difficult to discern whether a correlation to [Fe/H] exists for Ba with so few stars.

5200 0.20

0.10 5000 0.00

[Zr/Fe] -0.10 4800 -0.20 4600 -0.40 -0.20 0.00 0.20 0.40

4400 (K) eff T 0.60 4200 0.40 0.20 4000

[Ba/Fe] 0.00 3800 -0.20 -0.40 -0.20 0.00 0.20 0.40 3600 [Fe/H]

Figure 3.15. Zr and Ba plotted as a function of metallicity, where again the dashed line represents solar values. Zr appears flat as a function of metallicity. A correlation to [Fe/H] for Ba may exist, but difficult to discern for a small sample of stars

A larger sample of open cluster stars may better reveal any correlation to temperature or metallicity that might exist for barium.

70 Chapter 4

A Large Uniform Cluster Sample

4.1 The First OCCAM Neutron-Capture Sample of

Open Clusters

Using follow-up observations of the APOGEE-based OCCAM survey, I have derived abundances for a uniform sample of 30 open clusters. In this chapter, I present the remaining sample of 24 less well-studied open clusters with measured abundances

4.1.1 Stellar Parameters

For many of these stars, stellar parameters and abundance determinations are unavailable in literature for comparison. Although I do not offer a full comparison to APOGEE derived abundances, I compare stellar parameters for the full sample in Tables 4.1. For all other tables in this section, I take advantage of the abbreviated APOGEE IDs defined in Table 3.1.

71 Table 4.1. Measured Stellar Parameter Comparison to APOGEE DR14 .

This Work APOGEE DR14

Star Teff log g Vmicro [Fe/H] Teff log g Vmicro [Fe/H]

(K) (dex) (km s−1) (dex) (K) (dex) (km s−1) (dex)

ASCC 14 2M05215118+3520310 4810 3.00 1.35 0.15 4812 3.12 1.20 0.21 ASCC 26 2M06503895+0714131 3875 1.10 1.70 -0.09 4009 1.03 1.87 -0.30 2M06502079+0715161 4555 2.03 1.52 -0.08 4556 2.23 1.56 -0.00 Berkeley 9 2M03323179+5239538 4925 2.90 1.50 -0.07 5020 2.98 1.34 -0.05 Berkeley 17 2M05202386+3037219 4100 2.00 1.70 -0.19 4279 1.81 1.50 -0.12 2M05203799+3034414 4140 1.65 1.47 -0.19 4302 1.89 1.51 -0.08 Berkeley 19 2M05240941+2937217 4475 1.90 1.50 -0.17 4514 2.08 1.70 -0.23 2M05240726+2934330 4820 2.50 1.30 -0.26 4882 2.44 1.56 -0.29 Berkeley 31 2M06573113+0816069 3850 0.80 1.86 -0.27 3890 0.88 2.02 -0.33 Berkley 53 2M20555767+5103206 4860 2.40 1.54 -0.13 4945 2.47 1.57 -0.07 Berkeley 85 2M20184497+3744174 3980 1.40 1.90 -0.09 FLa FLa FLa FLa 2M20190397+3745002 3880 1.40 2.05 -0.12 FLa FLa FLa FLa 2M20185346+3745129 3775 1.10 2.00 -0.08 FLa FLa FLa FLa Berkley 91 2M21105678+4831562 4760 2.30 1.70 -0.40 4845 2.30 1.53 -0.54 Collinder 106 2M06381107+0552571 4330 1.80 1.80 -0.17 4422 1.92 1.66 -0.02 2M06381152+0537411 4870 2.40 1.58 -0.34 5015 2.51 1.55 -0.26 FSR 0498 2M00291126+6225307 4570 2.10 1.47 -0.30 4624 2.29 1.49 -0.23 IC 1369 2M21120996+4744158 4905 2.40 1.90 -0.03 4998 2.47 1.76 -0.03 2M21115265+4744571 4930 2.40 1.50 -0.03 4978 2.42 1.84 -0.01 2M21120610+4745175 4765 2.40 1.65 0.05 4816 2.51 1.59 0.17 2M21121345+4745256 4920 2.40 2.00 0.02 4995 2.31 2.15 -0.04 King 5 2M03142548+5247355 4240 1.70 1.88 -0.13 4306 1.64 1.72 -0.07 2M03140915+5237511 5095 2.75 1.53 -0.15 5150 2.64 1.55 -0.12 Melotte 71 2M07373529-1202311 5100 2.80 1.35 -0.25 5158 2.75 0.72 -0.14 2M07373589-1205094 4985 2.70 1.25 -0.28 5080 2.64 1.22 -0.10 NGC 1798 Continued on Next Page. . .

72 Table 4.1 – Continued

This Work APOGEE DR14

Star Teff log g Vmicro [Fe/H] Teff log g Vmicro [Fe/H]

(K) (dex) (km s−1) (dex) (K) (dex) (km s−1) (dex)

2M05113666+4741482 4720 2.40 1.70 -0.18 4749 2.28 1.63 -0.17 2M05114006+4739238 4700 2.20 1.60 -0.14 4765 2.28 1.60 -0.17 NGC 1817 2M05134177+1643290 5025 2.70 1.50 -0.09 5151 2.66 1.46 -0.08 2M05123284+1628255 4650 2.00 1.60 -0.18 ...... NGC 1896 2M05261353+2920425 4975 2.50 1.60 -0.09 4980 2.56 1.50 -0.12 2M05255232+2911527 5075 2.60 1.50 -0.07 5107 2.70 1.42 -0.03 2M05261358+2916170 4950 2.40 1.60 -0.12 5011 2.41 1.66 -0.37 NGC 1912 2M05281800+3545224 4950 2.50 1.65 -0.15 4971 2.81 1.46 0.02 NGC 2240 2M06332299+3513447 4845 2.30 1.40 -0.02 4834 2.53 1.48 0.08 NGC 2355 2M07170625+1341307 4755 3.10 1.29 0.13 4754 3.00 1.18 0.22 NGC 6705 2M18510399-0620414 4700 2.00 1.72 0.05 4796 2.40 1.82 0.20 2M18505944-0612435 4865 2.65 1.70 0.03 4952 2.66 1.05 0.14 2M18510092-0614564 4805 2.45 1.50 0.01 4853 2.51 1.36 0.15 2M18510786-0617119 4780 2.55 1.55 0.07 4861 2.59 1.33 0.13 NGC 6811 2M19365712+4622425 5115 2.95 1.28 -0.04 5108 2.74 1.20 -0.03 2M19373462+4624098 4990 2.40 1.47 0.00 4931 2.56 1.46 0.02 NGC 7062 2M21232731+4622136 5050 2.60 1.65 0.02 5087 2.60 1.59 0.04 Ruprecht 24 2M07314150-1244132 4925 3.00 1.40 -0.40 5043 3.35 1.13 -0.25 2M07315152-1247237 4830 2.30 1.40 -0.32 4877 2.43 1.70 -0.29

a FL = First Light APOGEE commissioning plate, calibrated abundances not available for APOGEE com-

missioning data in DR14.

Radial velocities for each star in this sample, with combined internal and external errors, are compared to APOGEE and presented in Table 4.2.

73 Table 4.2. Measured Paramters and Radial velocities.

Star Teff log g Vmicro [Fe/H] σ[F e/H] Vr σVr APO Vr APO Vscat (K) (dex) (km s−1) (dex) (dex) (km s−1) (km s−1) (km s−1) (km s−1)

ASCC 14 2M05215118+3520310 4810 3.00 1.35 0.15 0.06 −21.58 0.08 −22.23 0.11 ASCC 26 2M06503895+0714131 3875 1.10 1.70 −0.09 0.06 13.17 1.00 11.68 0.03 2M06502079+0715161 4555 2.03 1.52 −0.08 0.06 4.70 0.86 3.75 8.24 Berkeley 9 2M03323179+5239538 4925 2.90 1.50 −0.07 0.06 −18.01 0.09 −18.04 0.03 Berkeley 17 2M05202386+3037219 4100 2.00 1.70 −0.19 0.05 −67.07 0.18 −68.69 1.65 2M05203799+3034414 4140 1.65 1.47 −0.19 0.06 −72.45 0.06 −73.10 0.15 Berkeley 19 2M05240941+2937217 4475 1.90 1.50 −0.17 0.08 18.20 0.08 17.44 0.08 2M05240726+2934330 4820 2.50 1.30 −0.26 0.07 17.21 0.05 17.37 0.02 Berkeley 31 2M06573113+0816069 3850 0.80 1.86 −0.27 0.04 54.41 0.12 59.43 0.31 Berkley 53 2M20555767+5103206 4860 2.40 1.54 −0.13 0.05 −33.83 0.08 −34.24 0.06 Berkeley 85 2M20184497+3744174 3980 1.40 1.90 −0.09 0.05 −36.17 0.09 −34.91 0.05 2M20190397+3745002 3880 1.40 2.05 −0.12 0.04 −36.48 0.09 −35.49 0.07 2M20185346+3745129 3775 1.10 2.00 −0.08 0.06 −33.46 0.09 −33.97 0.17 Berkley 91 2M21105678+4831562 4760 2.30 1.70 −0.40 0.07 −19.42 1.55 −18.99 0.12 Collinder 106 2M06381107+0552571 4330 1.80 1.80 −0.17 0.07 35.79 0.08 34.74 0.02 2M06381152+0537411 4870 2.40 1.58 −0.34 0.07 36.73 0.18 36.81 0.08 FSR 0498 2M00291126+6225307 4570 2.10 1.47 −0.30 0.07 −10.21 0.07 −9.42 0.08 IC 1369 2M21120996+4744158 4905 2.40 1.90 −0.03 0.05 −48.40 0.08 −48.79 0.05 2M21115265+4744571 4930 2.40 1.50 −0.03 0.05 −48.72 0.09 −48.94 0.11 2M21120610+4745175 4765 2.40 1.65 0.05 0.06 −50.81 0.09 −50.53 0.07 2M21121345+4745256 4920 2.40 2.00 0.02 0.06 −48.22 0.15 −48.97 0.17 King 5 2M03142548+5247355 4240 1.70 1.88 −0.13 0.07 −43.30 0.04 −43.45 0.08 2M03140915+5237511 5095 2.75 1.53 −0.15 0.07 −45.48 0.16 −42.59 0.13 Melotte 71 2M07373529−1202311 5100 2.80 1.35 −0.25 0.07 50.80 0.06 51.36 0.13 2M07373589−1205094 4985 2.70 1.25 −0.28 0.07 50.81 0.09 50.64 0.08 NGC 1798 2M05113666+4741482 4720 2.40 1.70 −0.18 0.08 3.11 0.11 2.88 0.06 2M05114006+4739238 4700 2.20 1.60 −0.14 0.09 2.73 0.06 2.42 0.02 Continued on Next Page. . .

74 Table 4.2 – Continued

Star Teff log g Vmicro [Fe/H] σ[F e/H] Vr σVr APO Vr APO Vscat (K) (dex) (km s−1) (dex) (dex) (km s−1) (km s−1) (km s−1) (km s−1)

NGC 1817 2M05134177+1643290 5025 2.70 1.50 −0.09 0.07 65.74 0.08 66.20 0.04 2M05123284+1628255 4650 2.00 1.60 −0.18 0.06 65.92 0.06 NAPO 0.00 NGC 1896 2M05261353+2920425 4975 2.50 1.60 −0.09 0.07 5.86 0.12 5.82 0.06 2M05255232+2911527 5075 2.60 1.50 −0.07 0.07 4.97 0.09 4.51 0.07 2M05261358+2916170 4950 2.40 1.60 −0.12 0.01 0.43 1.03 0.80 0.14 NGC 1912 2M05281800+3545224 4950 2.50 1.65 −0.15 0.08 6.31 0.14 6.01 0.19 NGC 2240 2M06332299+3513447 4845 2.30 1.40 −0.02 0.06 −10.80 0.08 −10.82 0.03 NGC 2355 2M07170625+1341307 4755 3.10 1.29 0.13 0.05 34.34 0.11 34.82 0.09 NGC 6705 2M18510399−0620414 4700 2.00 1.72 0.05 0.07 34.04 0.09 34.86 0.00 2M18505944−0612435 4865 2.65 1.70 −0.03 0.07 35.44 0.09 34.87 0.01 2M18510092−0614564 4805 2.45 1.50 0.01 0.06 35.30 0.19 34.97 0.54 2M18510786−0617119 4780 2.55 1.55 0.07 0.05 37.35 1.54 34.71 0.01 NGC 6811 2M19365712+4622425 5115 2.95 1.28 0.04 0.06 8.37 0.09 7.91 0.00 2M19373462+4624098 4990 2.40 1.47 −0.00 0.06 8.60 0.11 7.66 0.00 NGC 7062 2M21232731+4622136 5050 2.60 1.65 0.02 0.08 −22.27 0.04 −22.26 0.14 Ruprecht 24 2M07314150−1244132 4925 3.00 1.40 −0.40 0.09 −58.05 1.54 −57.49 0.27 2M07315152−1247237 4830 2.30 1.40 −0.32 0.09 −51.46 0.11 −51.74 0.11

4.1.2 Berkeley 85

Even though Berkeley 85 has been listed as a possible moving group in a study by Bica et al.(2002), I include it here to present a method commonly used by stellar spectro- scopics for determining initial estimates for Teff and log g. We offer a color-temperature relationship and standard relation for identifying an approximate temperature and sur- face gravity for stars.

75 The bright stars in Berkeley 85 were utilized for commissioning the IR spectrograph

currently employed by APOGEE at the Apache Point Observatory in New Mexico. As

such, no stellar parameters from the APOGEE database were available as initial estimates

for our optical follow-up. We employed the color-temperature relationship described by

Alonso et al.(1999), specifically their equation 8, based on the infrared flux method

(Blackwell & Shallis 1977). The color transformations depend on V magnitudes (provided

by Zacharias et al. 2004) for V −K temperature estimates. For the full set of color transformations used, see Johnson et al.(2005), and references therein.

Surface gravities were calculated using the standard relation,

log(g) = 0.40(Mbol. − Mbol. + log(g ) + 4(log(T/T )) + log(M/M ), (4.1)

and assumed a stellar mass of of 0.8 M . The absolute bolometric magnitudes (Mbol.) were determined by applying the V -band corrections from Alonso et al. (1999; their equations 17 and 18) to the absolute V -band magnitudes estimated from the distance modulus (m−M)V = 11.23 and E(B − V ) = 0.77.

Our final temperatures and surface gravities are in excellent agreement with the initial estimates determined by using this method. For Berkeley 85 2M20184497+3744174 the

final Teff and log g were 3980 K and 1.40 dex, with initial estimates of 3962 K and 1.37 dex. For the two other stars, 2M20190397+3745002 and 2M20185346+3745129, initial estimates were Teff = 3836 K, log g =1.54 dex and Teff = 3801 K, log g =1.15 dex, respectively. Our final Teff = 3880 K, log g =1.40 dex for 2M20190397+3745002 and

Teff = 3775 K, log g =1.10 dex for 2M20185346+3745129 are in excellent agreement.

76 Abundances derived using these stellar parameters can be found in Table 4.3 and in

AppendixB with the rest of the cluster sample. Table 4.3. Measured Abundances for Berkeley 85.

Ratio 4147 5002 5129

[FeI/H] −0.09±0.05 −0.12±0.05 −0.08±0.06 [FeII/H] −0.09±0.05 −0.09±0.05 −0.06±0.07 [Na/Fe] 0.25±0.06 0.21±0.05 0.05±0.07 [Mg/Fe] −0.19±0.03 0.03±0.00 −0.08±0.00 [Al/Fe] −0.04±0.01 0.04±0.07 −0.07±0.10 [Si/Fe] 0.15±0.05 0.11±0.02 0.11±0.07 [Ca/Fe] −0.16±0.07 −0.19±0.10 −0.23±0.10 [ScII/Fe] −0.02±0.05 −0.11±0.04 −0.04±0.07 [Ti/Fe] −0.02±0.07 0.01±0.03 −0.12±0.05 [TiII/Fe] 0.02±0.03 0.08±0.06 0.01±0.06 [V/Fe] −0.16±0.05 −0.22±0.06 0.01±0.11 [Ni/Fe] 0.00±0.06 0.00±0.03 −0.04±0.08 [Zr/Fe] −0.20±0.06 −0.39±0.05 −0.30±0.09 [BaII/Fe] 0.14±0.00 0.19±0.00 −0.02±0.00

4.2 Summary

In Figure 4.1 I include the abundances determined in this study with the boxplot taken from 20 studies of 37 open clusters presented in Figure 2.1 of chapter2. The median abundance values for many elements have moved to more evenly reflect a true median, but the spread for most of the elements remains unchanged, or has expanded.

77 1.0

0.8

0.6

0.4

0.2 [X/Fe] 0.0

-0.2

-0.4

-0.6 [Fe/H] O Na Mg Al Si Ca Sc Ti V Mn Ni Y Zr Ba La Ce Nd Eu

Figure 4.1. Boxplot from Chapter2 is presented with the addition of abundances from this study. The middle line of each box indicates the median abundance value, and the upper and lower box boundaries represent the third and first quartiles (75th and 25th percentile) of the data, respectively. The vertical lines represent the full range of abundance values; the dashed line indicates solar.

78 We can now assume the large spreads have physical origins. With the smaller sample of stars, we saw a hint of temperature dependence for Zr abundances. With a larger

0.6 10.0 0.4 0.2 0.0 9.5 [Zr/Fe] -0.2 -0.4 4000 4200 4400 4600 4800 5000 9.0

Teff (K)

0.6 8.5 log(Age/yr) 0.4 0.2 0.0

[Zr/Fe] 8.0 -0.2 -0.4

-0.40 -0.20 0.00 0.20 7.5 [Fe/H]

Figure 4.2. Zirconium abundances for the full sample. The temperature dependence is removed with a larger sample of stars. The metallicity remains flat for all Zr abundances.

sample of stars we see only a small fraction of stars that show an over-abundances with respect to temperature. Other warm stars in this sample show an under-abundance as well as the relatively cool stars. If the super solar stars are removes from the data set, the trend is also removed and abundances are flat across both temperature and metallicity.

79 As with zirconium, the full sample of stars in this study reveal no dependence for either temperature or metallicity for barium abundances in open clusters.

0.8 10.0 0.6 0.4 0.2 9.5 [Ba/Fe] 0.0 -0.2 4000 4200 4400 4600 4800 5000 9.0

Teff (K) 0.8 8.5 log(Age/yr) 0.6 0.4 0.2 8.0 [Ba/Fe] 0.0 -0.2 -0.40 -0.20 0.00 0.20 7.5 [Fe/H]

Figure 4.3. Barium abundances for the full sample. No trends are aparent for Ba II abundances for either temperature or metallicity.

80 We return to the study by D’Orazi et al.(2009) and plot our clusters as a function of age and find a similar trend of decreasing Ba abundance with increasing age. Abundances appear to peak at log(Age) ∼8.5.

0.80 This study, possible moving group This study 0.60

0.40

[Ba/Fe] 0.20

0.00

-0.20 7.50 8.00 8.50 9.00 9.50 10.00 log(Age/yr) Figure 4.4. Clusters from this study are plotted as a function of age. Bona fide open clusters are shown as purple filled circles, while possible moving groups shown as filled red circles.

81 Clusters from this study are plotted along with the literature compilation by D‘Orazi in Figure 4.5 showing a strong trend for the previously mentioned age-metallicity relation.

However, in a very real sense, this now becomes another literature compilation inheriting the same weaknesses— specifically, different methodologies for abundance derivation, and with different stellar types.

0.80 This study, possible moving group giants (D’Orazi 2009) dwarfs (D’Orazi 2009) 0.60 This study

0.40

[Ba/Fe] 0.20

0.00

-0.20 7.50 8.00 8.50 9.00 9.50 10.00 log(Age/yr) Figure 4.5. Over-plot of clusters from this study to the literature compilation from the D‘Orazi study, we begin to see a trend emerge. Giants from the D‘Orazi study are shown as open squares, dwarfs as filled black squares. Plotting designations for this study remain the same as Figure 4.4.

82 By removing the dwarfs from the D’Orazi study and compare to our sample of giants, a clearer trend develops for the age-metallicity relationship.

0.80 This study, possible moving group giants (D’Orazi 2009) This study 0.60

0.40

[Ba/Fe] 0.20

0.00

-0.20 7.50 8.00 8.50 9.00 9.50 10.00 log(Age/yr) Figure 4.6. Removing all but giant stars from the sample from the D‘Orazi sample, a clear trend of decreasing [Ba/Fe] abundance with increasing age can be seen from both studies beginning just after log(Age) ∼8.5.

Recall, this study has severely restricted its use of Ba absorption lines to permit only measurements less than 150 mA.˚ In most cases, this limits our measurements to a single

Ba line for each star, λ5853 A.˚ Even so, in this study’s view, it returns the most reliable abundance measurement than the other transition lines available in our wavelength range.

83 I can then confirm the trend found by D‘Orazi for decreasing [Ba/Fe] abundance from

9.0 < log (age) < 10; however I find that the trend ends around log (age) ∼ 9.0 and that either I have a constant of [Ba/Fe] ∼ 0.5 or reversing slope for ages younger than a Gyr.

Additional younger clusters are needed to definitively confirm the [Ba/Fe] trend between

6.0 < log (age) < 9.0.

4.3 Discussion

Astronomers have long used star clusters as empirical testbeds for the purposes of under- standing both the kinematic properties of the Galaxy as well as its chemical evolution.

These bound stars share a common formation history and, as a result, share a common age, distance from us, relative velocity and spatial relationship on the sky. Open clusters are loosely bound, relatively young, star clusters that reside almost exclusively in the plane of the Milky Way making them particularly useful systems, in aggregate, for inves- tigating Galactic dynamical and chemical evolutionary processes. Having been formed from the same well-mixed giant molecular cloud (GMC), they are also thought to be fairly homogeneous in their chemical composition (e.g De Silva et al. 2006; 2007a). For many elements, the observed star-to-star abundance variations for constituent members within open clusters show a remarkably small spread, ∼0.01 to 0.05 dex (Pancino et al.

2010, De Silva et al. 2011, Ting et al. 2012, Reddy et al. 2012; 2013b, De Silva et al. 2013,

Bertran de Lis et al. 2016). These common attributes acquired from a shared formation history allow astronomers to differentiate stars that may belong to a particular cluster from those that share the same field of view (O’Connell et al. 2016).

84 Even though the vast majority of stars, if not all, are born in clusters (Lada & Lada

2003), the array of stars we see in the night sky are, for the most part, either single stars or binary systems suggesting some disruption of clusters must have occurred on a relatively short timescale (Lada & Lada 2003, Bastian et al. 2005, Mengel et al. 2005). The Galactic disk is a dynamic environment where close encounters or interactions with other clusters, stars in the field or GMCs can disrupt a cluster, causing it to lose member stars to the disk. Perturbations due to spiral arm rotation and the central Galactic bar cause orbital variations that may disperse cluster members, as well as orbital resonances that can trap these dispersed member stars (Dehnen 1998). Both internal star formation dynamics (see

Lada 2010, for a review) and external Galactic perturbations can eject cluster member stars, where they migrate to the disk as unbound stars or moving groups, and become part of the general stellar disk population. These are the field stars. As an isolated star, our own Sun is considered a field star.

While Galactic orbital elements may change for these stars, their chemical composi- tion is preserved, as well as the small star-to-star elemental scatter (see De Silva et al.

2007b, Bubar & King 2010). High resolution, high SNR spectroscopic abundance analy- ses can reveal the chemical signatures necessary for identifying such stellar associations for the purpose of “chemically tagging,” i.e. tagging stars with no known or otherwise discernible associations, to either the thick or thin disk of the Galaxy (Freeman & Bland-

Hawthorn 2002, Hogg et al. 2016, Kos et al. 2017). While beyond the scope of this study, chemically tagging stars to a coeval, or conatal, star cluster in the interest of reconstruct- ing the chemical evolution of the Galaxy has become an intriguing proposition for near

field cosmology.

85 4.3.1 Chemo-chronology

Dating stars by chemistry has been a matter of investigation in stellar spectroscopy for decades (Laird & Sneden 1996, Fischer 1998, Charbonnel & Vauclair 1999, Charbonnel

1999, Sneden et al. 2001, Cowan & Sneden 2006, Li et al. 2012). Methods for detecting an age-metallicity relation are generally aimed at detecting ages for the large population of field stars in the Galaxy, for which we currently have no means for assessing their ages. Lithium depletion, or lithium burning, in stars has been the favored tracer of age

(see Randich 2009, Randich et al. 2009). The high temperatures necessary for hydrogen fusion rapidly depletes the primordial lithium in stars. Through the collision of a proton with7Li, two 4He nuclei are created. Yet, Li-depletion ages disagree with isochronal ages, and stellar rotation may be a factor in mixing mechanisms that prevent Li from being detected on the surface of a star (Brott et al. 2011).

Both Sneden et al.(2001) and Freeman & Bland-Hawthorn(2002) suggest that n- capture elements, specifically the r-process elements, offer promising possibilities for chemically tagging stars to age in very metal-poor stars (−3.0 ≤ [Fe/H] ≤ -2.0). These ultra-metal-poor stars, however, generally reside in the Galactic halo and few have been identified or analyzed for chemical abundances in detail. Further, the necessary analysis for many r-process elements requires not only high resolution spectra, but SNRs suffi- ciently high enough, ∼1000, to eliminate the inevitable blending of spectral lines (Kurucz

1991, Kurucz & Bell 1995).

86 4.3.2 The Problem with Barium

In relatively cool stellar atmospheres, Ba II has only five strong transitions from lower energy levels in the optical regime. At least two of these lines tend either to be saturated

(λ6496 A)˚ or terminally blended with iron (λ6141 A).˚ Other transitions, with higher excitation potentials, are very weak. Significant hyperfine and isotope splitting exists for all Ba II lines, both with high and low energy levels. Barium also has five stable isotopes, and their relative abundances are synthesized by both the s- and r-processes, which can play a role in abundance derivations (Magain 1995). Ba II is also highly sensitive to the microturbulent velocity chosen as a stellar parameter for analysis (see Table 3.15).

The importance of an easily observable, easily measurable element, or group of el- ements, in the metal-rich Galactic disk which would constrain ages for field stars, e.g. as in dissolved open clusters, cannot be overstated. Reconstructing the evolution of our Galaxy requires well constrained ages for not only cluster stars, but also for stars that have been dislocated from their natal cluster by dynamical Galactic processes. The

[Ba/Fe] relationship with age presented in this study looks promising. However, more homogeneous studies involving more clusters, including a more statistically significant number of cluster members, is needed to fully investigate the relationship.

This confirmed [Ba/Fe] trend measurement may provide a key stellar evolutionary time scale constraint for Galaxy evolution modelers in the future, but most models do not yet have the chemical abundance patterns detailed for individual elements. I find that [Ba/Fe] is a key element that should be included in more advanced models, but also caution other observers about a careful treatment of the Ba lines used.

87 Chapter 5

Moving Forward

What do astronomers do? We turn

starlight into numbers, and then we talk

about those numbers for about twenty

years. Then those numbers change, and

we talk about why they changed for

another twenty years.

John W. Kuehne,

McDonald Observatory, 2014

It is well known that no age-metallicity relation exists in Galactic open clusters.

There are more than 2000 open clusters known to exist in the Galaxy, of a likely sample of ∼10,000 open clusters. Some think many more may exist. In the larger view, this is a slim sample. More clusters samples, with larger member numbers, should be investigated to confirm or deny the relationship seen in literature compilations, or the one seen with a relatively small sample presented here. By relatively, I mean relative to the to the number

88 of open clusters in the Galaxy. Ideally, I would continue investigating this development, but this was never my intended goal.

To better explore beyond the Galactic disk, I will expand my skill set by analyzing different types of stars, in different components of the Galaxy– outside the Galactic plane where extremely metal-poor stars can be found, or toward the center of the Galaxy, obscured in the optical bands by a concentration of dust. Additional r-process elements can be included, as well as important elements such as oxygen, carbon and nitrogen.

I’m going to need a bigger telescope.

In the age of large sky surveys, where the means and skills exist for determining abun- dances for 100,000’s of stars, one might tend to think that by-hand analyses of stars for chemical abundances would be a skill set in decline for demand. On the contrary. Large sky surveys demand a calibrating influence for the pipelines responsible for producing accurate results on such a large scale. By-hand analyses of stars for chemical abundances offers just that calibrating influence. There is no substitute for first hand inspection of the data, reliable analysis techniques and interpretation of results that may differ from large scale analyses.

A promising postdoctoral position and fellowship, supported by the Universidad de

Concepci`onin Chile, will allow me to investigate some of the oldest, metal-poor star clusters in the Galactic bulge and halo. The bulge Cluster APOgee Survey (CAPOS), in conjunction with APOGEE-2 South, will investigate the chemical composition of star clusters, both globular and open, and determine ages for some of the oldest existing stellar members of the Galaxy. This will strongly constrain the age of the Galaxy and put it in the context of the timeline of formation after the Big Bang. Under the direction of Doug

89 Geisler I would continue to have access to all APOGEE data. Additionally, I would be encouraged, and expected, to utilize my particular skill set for analyzing abundances for stars observed by the project outside of APOGEE. This project aligns very well with my project detailed in this thesis, and would extend my research to star clusters in multiple components of the Galaxy. The opportunity may also afford a further investigation into the [Ba/Fe] v. age relationship with star clusters at larger Galactic distances.

90 Appendix A

Cluster Line lists

In this section we present the full line lists and equivalent width measurements for all stars in the calibration clusters, M67, NGC 188, NGC 2420, NGC 6819, NGC 6791, and

NGC 7789.

91 Table A.1. M67 Line list.

λ(A)˚ Element E.P. (eV) log(gf) 0076 1230 3520 4263 4264 5049 6061 6425 7168 9493 5417.037 FeI 4.415 -1.68 ...... 56.20 ...... 51.00 ...... 5441.339 FeI 4.312 -1.73 ...... 62.10 ...... 5445.042 FeI 4.386 -0.21 ...... 5466.396 FeI 4.371 -0.63 ...... 110.00 ...... 106.90 . . . 107.60 112.20 . . . 5466.987 FeI 3.573 -2.23 ...... 74.10 ...... 76.10 . . . 74.10 86.20 . . . 5505.881 FeI 4.412 -1.20 ...... 5522.446 FeI 4.209 -1.55 ...... 75.30 ...... 5557.897 FeI 3.111 -3.71 ...... 5576.089 FeI 3.430 -1.00 ...... 149.30 ...... 148.80 . . . 155.80 ...... 5679.024 FeI 4.652 -0.92 73.30 72.70 75.20 83.60 80.30 75.30 76.30 82.00 79.30 69.40 5731.762 FeI 4.256 -1.30 77.40 78.80 92.20 88.50 79.80 83.00 80.40 90.90 88.20 70.90 5752.032 FeI 4.548 -0.92 . . . 81.40 85.30 84.40 81.90 80.60 84.70 89.50 92.40 66.90 5760.344 FeI 3.642 -2.49 53.30 58.60 65.70 63.70 62.20 66.00 53.00 64.90 73.70 45.60 5775.081 FeI 4.220 -1.30 72.60 85.70 86.50 82.80 82.10 83.70 80.90 82.60 89.30 69.40 5778.453 FeI 2.588 -3.59 59.10 61.50 71.40 69.60 71.20 75.20 69.50 72.60 84.90 44.60 5793.915 FeI 4.220 -1.70 . . . 61.10 73.50 68.90 69.70 65.90 58.60 68.30 72.40 52.20 5806.726 FeI 4.607 -1.05 64.20 75.00 83.90 75.60 74.00 78.30 74.10 81.30 78.50 60.60 5855.076 FeI 4.607 -1.76 30.80 38.40 42.30 39.40 42.30 44.70 35.80 45.00 48.20 29.80 5856.088 FeI 4.294 -1.64 51.20 62.20 70.40 68.60 64.00 64.80 57.20 72.90 73.10 49.20 5916.247 FeI 2.458 -2.99 102.80 104.70 120.10 112.40 110.30 106.70 109.00 116.30 127.40 83.20 5927.789 FeI 4.652 -1.09 58.40 65.60 69.50 71.10 73.10 65.40 66.70 68.40 70.80 . . . 5929.677 FeI 4.548 -1.41 . . . 59.70 62.10 64.20 62.20 61.60 52.30 67.10 65.70 48.30 5934.655 FeI 3.928 -1.17 98.10 103.70 111.10 114.40 111.60 117.90 103.60 116.00 117.60 89.40 5987.065 FeI 4.795 -0.43 91.30 90.40 . . . 101.60 94.50 92.40 93.40 98.90 96.10 82.80 6027.051 FeI 4.076 -1.21 86.70 93.90 98.60 99.70 95.80 102.80 ...... 103.70 . . . 6056.004 FeI 4.733 -0.46 . . . 92.30 98.00 102.60 97.90 97.80 94.90 95.20 102.10 88.10 6079.008 FeI 4.652 -1.12 63.40 66.30 73.50 77.20 69.40 71.20 69.20 71.90 74.50 . . . 6093.643 FeI 4.607 -1.50 . . . 51.60 57.80 59.80 52.80 52.90 52.80 52.50 61.10 . . . 6096.664 FeI 3.984 -1.93 59.30 64.80 75.60 69.10 68.60 77.90 61.40 74.10 77.40 47.80 6151.617 FeI 2.176 -3.37 93.00 101.10 . . . 108.20 104.30 106.20 99.50 106.70 126.30 . . . 6165.360 FeI 4.143 -1.55 65.80 72.90 77.80 83.20 74.10 77.50 73.20 76.20 87.60 60.80 6180.203 FeI 2.727 -2.78 95.20 102.80 116.40 110.10 ...... 107.30 118.60 127.10 . . . 6187.989 FeI 3.943 -1.72 71.10 75.30 83.30 87.40 84.70 84.70 72.60 79.50 89.90 60.50

92 Continued on Next Page. . . Table A.1 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 0076 1230 3520 4263 4264 5049 6061 6425 7168 9493 6200.313 FeI 2.608 -2.44 114.30 118.90 129.60 128.80 126.00 124.80 124.10 123.00 146.90 100.50 6229.226 FeI 2.845 -2.97 75.80 87.10 94.50 87.50 85.50 90.20 82.30 87.10 102.20 60.40 6232.641 FeI 3.654 -1.24 113.60 121.60 134.50 133.20 127.70 126.40 124.60 121.70 137.20 102.20 6240.645 FeI 2.223 -3.38 97.50 103.50 116.40 106.60 107.10 108.50 102.50 106.90 127.00 . . . 6246.318 FeI 3.602 -0.77 144.40 ...... 156.40 ...... 163.00 . . . 132.50 6270.223 FeI 2.858 -2.71 83.40 90.90 101.50 107.40 95.90 103.30 99.10 96.90 112.70 75.20 6301.501 FeI 3.654 -0.71 ...... 156.60 154.80 ...... 138.80 6336.824 FeI 3.686 -1.05 124.10 133.40 146.00 142.20 . . . 139.30 137.50 137.90 146.90 . . . 6393.600 FeI 2.433 -1.62 ...... 6475.624 FeI 2.559 -2.94 95.20 ...... 108.70 . . . 114.60 ...... 128.80 . . . 6481.869 FeI 2.279 -3.01 107.10 . . . 131.40 ...... 119.80 . . . 122.90 139.60 94.40 6494.980 FeI 2.404 -1.27 ...... 6498.940 FeI 0.958 -4.69 110.50 114.20 128.50 130.50 118.30 124.30 122.00 131.40 150.60 . . . 6533.928 FeI 4.558 -1.43 52.70 56.70 . . . 60.20 58.10 . . . 58.60 . . . 59.80 49.40 6546.238 FeI 2.758 -1.65 ...... 6569.214 FeI 4.733 -0.42 93.80 95.90 110.90 108.40 106.50 108.30 97.70 . . . 108.90 88.30 6575.016 FeI 2.588 -2.82 . . . 111.30 ...... 120.70 112.70 118.20 134.30 90.10 6592.913 FeI 2.727 -1.60 158.80 ...... 6593.870 FeI 2.433 -2.42 137.20 139.50 149.00 149.80 . . . 147.70 140.70 ...... 115.70 6597.559 FeI 4.795 -1.07 52.90 63.40 65.70 69.10 65.40 65.10 65.50 . . . 67.80 56.10 6608.025 FeI 2.279 -4.03 54.10 64.10 68.70 68.20 75.10 66.90 65.60 73.50 92.50 43.40 6609.110 FeI 2.559 -2.69 114.10 119.60 129.90 129.60 120.40 127.20 121.10 124.80 144.00 96.30 6609.678 FeI 0.990 -5.18 81.00 95.30 99.30 99.10 94.30 100.80 85.80 108.60 130.60 62.60 6627.545 FeI 4.548 -1.59 40.30 50.90 . . . 55.70 51.80 . . . 57.50 ...... 42.70 6646.931 FeI 2.608 -3.99 36.40 42.90 . . . 55.00 53.40 . . . 52.80 ...... 6677.985 FeI 2.692 -1.42 ...... 6703.566 FeI 2.758 -3.16 69.50 81.40 . . . 82.70 89.20 . . . 82.10 ...... 64.10 6705.102 FeI 4.607 -0.99 72.90 72.80 . . . 78.30 77.90 . . . 75.70 ...... 70.70 6710.318 FeI 1.485 -4.88 62.70 77.20 . . . 87.10 80.40 . . . 68.60 ...... 47.50 6713.743 FeI 4.795 -1.60 32.00 36.20 . . . 41.80 42.70 . . . 34.60 ...... 27.00 6725.357 FeI 4.103 -2.10 40.40 ...... 45.50 ...... 6726.666 FeI 4.607 -1.09 68.10 71.80 . . . 79.90 74.50 . . . 71.20 ...... 61.60 6733.150 FeI 4.638 -1.58 ...... 6793.258 FeI 4.076 -2.47 ...... 35.60 ...... 93 Continued on Next Page. . . Table A.1 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 0076 1230 3520 4263 4264 5049 6061 6425 7168 9493 6810.262 FeI 4.607 -1.12 . . . 71.80 . . . 79.60 78.80 ...... 64.60 6820.371 FeI 4.638 -1.32 54.90 62.30 . . . 62.90 66.30 . . . 61.40 ...... 47.00 6828.592 FeI 4.638 -0.92 71.20 81.40 . . . 80.40 82.90 . . . 84.80 ...... 70.30 6839.830 FeI 2.559 -3.45 74.10 80.60 . . . 84.70 87.90 . . . 75.60 ...... 61.20 6842.685 FeI 4.638 -1.32 ...... 6843.655 FeI 4.548 -0.93 ...... 5425.249 FeII 3.199 -3.36 ...... 48.40 ...... 54.40 ...... 41.50 . . . 6084.103 FeII 3.199 -3.81 21.60 28.00 29.90 31.70 30.80 30.80 ...... 24.00 29.30 6149.246 FeII 3.889 -2.73 29.10 38.30 44.60 42.80 34.70 44.70 . . . 42.00 29.00 . . . 6247.559 FeII 3.891 -2.35 44.20 52.60 . . . 64.20 52.00 63.70 47.00 . . . 43.70 56.20 6369.459 FeII 2.891 -4.25 17.20 22.10 27.60 30.90 24.70 29.00 23.40 25.40 23.30 . . . 6416.919 FeII 3.891 -2.75 . . . 34.60 44.90 44.90 38.40 46.20 ...... 42.60 6432.677 FeII 2.891 -3.71 33.80 40.70 46.50 53.20 48.10 51.60 41.90 50.00 40.30 47.30 6456.381 FeII 3.903 -2.09 54.60 58.80 ...... 66.40 67.30 60.80 . . . 53.10 68.30 6516.077 FeII 2.891 -3.37 ...... 62.70 ...... 62.10 ...... 6154.226 NaI 2.102 -1.55 74.30 79.90 80.50 85.80 81.40 80.80 81.60 71.90 108.10 45.70 6160.747 NaI 2.104 -1.26 96.20 101.90 100.50 97.50 103.50 96.70 101.90 99.30 120.30 70.80 5711.088 MgI 4.346 -1.83 125.70 131.20 144.30 135.70 134.70 131.10 132.70 139.40 146.30 122.40 6318.717 MgI 5.108 -1.73 77.90 85.70 91.20 86.00 89.40 91.00 84.80 86.50 91.10 77.00 6319.237 MgI 5.108 -1.95 63.10 ...... 6319.495 MgI 5.108 -2.43 36.10 ...... 6696.023 AlI 3.140 -1.35 82.50 87.70 . . . 92.70 85.80 . . . 83.40 ...... 70.20 6698.673 AlI 3.140 -1.87 47.60 51.10 . . . 50.20 58.70 . . . 54.40 ...... 33.70 5690.425 SiI 4.929 -1.87 54.60 61.70 64.90 63.60 58.60 60.80 61.30 62.80 58.60 . . . 5701.104 SiI 4.929 -2.05 41.40 42.40 51.20 59.40 . . . 55.30 50.00 . . . 50.70 47.10 5772.146 SiI 5.082 -1.75 62.30 . . . 67.60 ...... 67.70 ...... 61.30 . . . 5793.073 SiI 4.929 -2.06 . . . 52.50 ...... 54.30 49.90 56.80 ...... 6142.483 SiI 5.619 -1.30 . . . 41.70 . . . 45.90 42.00 51.90 ...... 54.70 6145.016 SiI 5.616 -1.31 42.00 42.20 . . . 47.20 43.90 56.00 42.90 ...... 49.60 6155.134 SiI 5.619 -0.65 84.30 81.90 93.40 89.50 79.50 86.70 87.60 92.70 77.10 . . . 6243.815 SiI 5.616 -1.24 48.20 ...... 51.80 . . . 54.00 6244.466 SiI 5.616 -1.09 . . . 56.30 65.30 59.20 52.20 61.00 66.90 . . . 60.30 . . . 6161.297 CaI 2.523 -1.02 . . . 116.60 120.20 118.80 114.90 112.40 111.70 119.40 140.80 . . . 6166.439 CaI 2.521 -0.90 110.00 113.10 . . . 116.70 124.00 117.70 ...... 144.60 99.30 94 Continued on Next Page. . . Table A.1 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 0076 1230 3520 4263 4264 5049 6061 6425 7168 9493 6169.042 CaI 2.523 -0.55 139.30 131.10 138.00 136.20 142.30 136.00 145.90 141.60 . . . 116.20 6169.563 CaI 2.526 -0.27 ...... 133.70 6455.598 CaI 2.523 -1.32 97.10 95.30 105.50 101.20 106.50 107.60 104.40 112.10 124.30 77.20 6471.662 CaI 2.526 -0.59 131.80 ...... 143.90 ...... 6493.781 CaI 2.521 0.14 ...... 6499.650 CaI 2.523 -0.59 . . . 134.20 139.60 138.10 . . . 136.20 143.40 ...... 6572.779 CaI 0.000 -4.29 ...... 71.50 5657.895 ScII 1.507 -0.60 84.80 92.40 96.50 98.10 89.20 99.80 83.90 99.90 106.50 78.50 5684.190 ScII 1.507 -1.07 ...... 67.20 74.30 74.10 64.40 . . . 79.60 . . . 6245.621 ScII 1.507 -1.02 65.60 71.20 79.00 74.60 70.90 70.20 58.10 71.60 92.20 56.30 6300.684 ScII 1.507 -1.90 . . . 26.80 . . . 21.80 21.20 25.40 ...... 30.70 . . . 6320.850 ScII 1.500 -1.82 . . . 23.40 31.10 26.10 28.50 29.80 . . . 34.30 36.50 22.90 6604.599 ScII 1.357 -1.31 55.70 68.70 . . . 73.20 64.40 77.20 60.50 74.60 84.50 . . . 5702.656 TiI 2.291 -0.65 40.50 44.50 39.90 42.00 56.10 52.30 52.80 48.90 ...... 5716.445 TiI 2.297 -0.78 . . . 34.20 28.60 32.50 38.90 36.50 34.90 . . . 73.60 17.80 6064.626 TiI 1.046 -1.94 53.50 63.80 60.00 56.80 75.50 60.80 66.60 64.20 108.40 31.90 6091.171 TiI 2.267 -0.32 . . . 64.40 . . . 66.80 72.40 . . . 63.90 ...... 44.70 6092.792 TiI 1.887 -1.38 36.20 ...... 39.90 50.80 33.00 ...... 66.40 . . . 6261.099 TiI 1.429 -0.53 ...... 124.40 ...... 81.10 6303.757 TiI 1.443 -1.59 46.70 50.40 44.50 48.20 60.90 55.70 57.00 . . . 102.90 . . . 6312.236 TiI 1.460 -1.55 45.10 52.30 . . . 56.50 62.10 . . . 58.30 56.10 94.80 26.90 6336.099 TiI 1.443 -2.01 ...... 31.10 . . . 44.90 40.80 . . . 37.30 ...... 6599.105 TiI 0.900 -2.08 56.80 69.30 60.80 69.30 75.10 66.50 67.10 . . . 123.30 35.80 5418.768 TiII 1.581 -2.13 ...... 85.80 . . . 79.30 92.90 . . . 6491.561 TiII 2.061 -1.94 49.30 59.20 61.20 68.10 67.90 . . . 58.60 ...... 52.80 6559.588 TiII 2.048 -2.18 42.10 44.10 43.80 51.80 50.90 56.10 50.30 52.00 60.70 . . . 6606.949 TiII 2.061 -2.79 ...... 23.40 ...... 20.40 6680.135 TiII 3.095 -1.89 ...... 6039.713 VI 1.064 -0.65 . . . 79.70 73.10 79.10 89.30 75.70 69.60 76.70 116.20 . . . 6081.443 VI 1.051 -0.61 71.50 75.70 84.70 90.30 92.10 87.30 74.60 . . . 125.30 39.60 6090.194 VI 1.081 -0.07 ...... 107.50 111.10 . . . 103.00 107.60 144.60 70.60 6119.525 VI 1.064 -0.36 92.70 99.20 96.60 97.80 . . . 86.60 95.60 ...... 57.70 6135.365 VI 1.051 -0.76 70.40 71.60 73.80 80.40 86.40 71.60 73.40 80.30 117.90 . . . 6274.652 VI 0.267 -1.70 67.40 77.20 75.10 76.30 89.00 76.70 75.40 76.10 134.90 27.30 95 Continued on Next Page. . . Table A.1 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 0076 1230 3520 4263 4264 5049 6061 6425 7168 9493 6285.160 VI 0.275 -1.54 . . . 81.90 83.70 85.30 93.80 88.80 82.70 . . . 142.00 44.20 6531.446 VI 1.218 -1.50 ...... 5760.828 NiI 4.105 -0.80 . . . 60.10 67.30 ...... 60.60 52.30 ...... 52.50 5805.213 NiI 4.167 -0.64 48.50 58.60 65.20 64.40 61.00 62.20 57.90 61.50 66.80 48.90 6176.812 NiI 4.088 -0.26 83.60 81.20 87.70 86.70 78.10 90.10 . . . 92.00 92.00 71.60 6204.604 NiI 4.088 -1.08 33.50 48.10 44.70 49.40 . . . 43.80 42.40 44.40 54.00 . . . 6223.984 NiI 4.105 -0.91 43.50 46.90 53.30 54.30 43.60 50.20 50.10 49.50 60.40 42.20 6378.260 NiI 4.153 -0.82 48.30 56.50 64.50 57.10 51.80 59.70 54.00 58.20 61.70 46.20 6127.475 ZrI 0.154 -1.06 20.00 23.70 26.80 31.30 36.70 22.80 25.60 . . . 74.20 9.10 6134.585 ZrI 0.000 -1.28 19.10 26.40 25.80 25.80 33.40 25.50 29.90 31.40 75.60 7.40 6140.535 ZrI 0.519 -1.41 ...... 24.20 . . . 6143.252 ZrI 0.071 -1.10 22.80 33.80 39.30 34.20 40.90 31.20 29.00 36.50 86.00 9.00 5853.668 BaII 0.604 -1.00 100.10 103.90 124.10 121.30 111.50 124.40 101.50 125.60 131.70 91.20 6141.713 BaII 0.704 -0.08 ...... 6496.897 BaII 0.604 -0.38 136.20 140.70 ...... 131.20 96 Table A.2. NGC 188 Line list.

λ(A)˚ Element E.P. (eV) log(gf) 3336 6329 6329 5417.037 FeI 4.415 -1.68 ...... 5441.339 FeI 4.312 -1.73 ...... 5445.042 FeI 4.386 -0.21 ...... 5466.396 FeI 4.371 -0.63 ...... 5466.987 FeI 3.573 -2.23 ...... 5505.881 FeI 4.412 -1.20 ...... 5522.446 FeI 4.209 -1.55 ...... 5557.897 FeI 3.111 -3.71 ...... 5576.089 FeI 3.430 -1.00 ...... 5679.024 FeI 4.652 -0.92 84.80 81.30 81.00 5731.762 FeI 4.256 -1.30 94.20 89.10 80.70 5752.032 FeI 4.548 -1.17 83.90 78.00 83.20 5760.344 FeI 3.642 -2.49 71.50 71.50 60.80 5775.081 FeI 4.220 -1.30 90.80 . . . 94.40 5778.453 FeI 2.588 -3.59 92.60 91.80 74.30 5793.915 FeI 4.220 -1.70 80.80 70.60 74.10 5806.726 FeI 4.607 -1.05 81.00 78.40 83.50 5855.076 FeI 4.607 -1.76 52.80 43.80 46.70 5856.088 FeI 4.294 -1.64 79.00 69.60 63.20 5916.247 FeI 2.458 -2.99 130.50 133.20 112.70 5927.789 FeI 4.652 -1.09 74.10 71.50 72.40 5929.677 FeI 4.548 -1.41 73.20 75.70 62.40 5934.655 FeI 3.928 -1.17 125.10 117.80 111.10 5987.065 FeI 4.795 -0.43 107.70 106.10 . . . 6027.051 FeI 4.076 -1.21 109.60 104.40 99.20 6056.004 FeI 4.733 -0.46 108.30 . . . 95.20 6079.008 FeI 4.652 -1.12 82.70 71.90 72.20 6093.643 FeI 4.607 -1.50 64.00 64.90 . . . 6096.664 FeI 3.984 -1.93 77.90 83.20 69.50 6151.617 FeI 2.176 -3.37 131.30 128.20 113.30 6165.360 FeI 4.143 -1.55 85.90 90.30 84.90 6180.203 FeI 2.727 -2.78 . . . 129.60 114.80 6187.989 FeI 3.943 -1.72 97.80 92.10 88.10 6200.313 FeI 2.608 -2.44 148.60 . . . 131.30 6229.226 FeI 2.845 -2.97 108.80 112.80 90.40 6232.641 FeI 3.654 -1.24 . . . 132.70 131.70 6240.645 FeI 2.223 -3.38 123.90 128.30 116.30 6246.318 FeI 3.602 -0.77 ...... 6270.223 FeI 2.858 -2.71 114.90 113.70 106.30 6301.501 FeI 3.654 -0.71 ...... 163.10 6336.824 FeI 3.686 -1.05 . . . 152.20 137.80 6393.600 FeI 2.433 -1.62 ...... 6475.624 FeI 2.559 -2.94 133.60 ...... 6481.869 FeI 2.279 -3.01 . . . 148.80 125.50 6494.980 FeI 2.404 -1.27 ...... 6498.940 FeI 0.958 -4.69 ...... 136.90 6533.928 FeI 4.558 -1.43 71.40 66.90 58.20 6546.238 FeI 2.758 -1.65 ...... 6569.214 FeI 4.733 -0.42 109.20 105.40 110.20 6575.016 FeI 2.588 -2.82 ...... Continued on Next Page. . .

97 Table A.2 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 3336 6329 6329 6592.913 FeI 2.727 -1.60 ...... 6593.870 FeI 2.433 -2.42 ...... 148.90 6597.559 FeI 4.795 -1.07 76.90 70.90 60.90 6608.025 FeI 2.279 -4.03 . . . 90.30 79.50 6609.110 FeI 2.559 -2.69 . . . 143.40 135.00 6609.678 FeI 0.990 -5.18 . . . 140.10 110.70 6627.545 FeI 4.548 -1.59 66.90 64.20 60.20 6646.931 FeI 2.608 -3.99 67.30 71.60 57.30 6677.985 FeI 2.692 -1.42 ...... 6703.566 FeI 2.758 -3.16 107.30 102.80 93.40 6705.102 FeI 4.607 -0.99 84.10 82.40 79.50 6710.318 FeI 1.485 -4.88 106.60 117.00 92.60 6713.743 FeI 4.795 -1.60 44.30 44.10 41.90 6725.357 FeI 4.103 -2.10 61.20 59.30 57.90 6726.666 FeI 4.607 -1.09 85.90 . . . 77.90 6733.150 FeI 4.638 -1.58 53.60 55.80 . . . 6793.258 FeI 4.076 -2.47 . . . 48.10 . . . 6810.262 FeI 4.607 -1.12 87.10 78.90 79.60 6820.371 FeI 4.638 -1.32 75.40 67.20 66.90 6828.592 FeI 4.638 -0.92 94.00 90.60 83.00 6839.830 FeI 2.559 -3.45 108.10 109.50 95.40 6842.685 FeI 4.638 -1.32 . . . 64.70 . . . 6843.655 FeI 4.548 -0.93 ...... 5425.249 FeII 3.199 -3.36 ...... 6084.103 FeII 3.199 -3.81 26.30 20.30 23.90 6149.246 FeII 3.889 -2.73 35.20 27.50 . . . 6247.559 FeII 3.891 -2.35 ...... 42.90 6369.459 FeII 2.891 -4.25 22.40 ...... 6416.919 FeII 3.891 -2.75 ...... 32.40 6432.677 FeII 2.891 -3.71 45.50 34.60 38.50 6456.381 FeII 3.903 -2.09 54.30 44.20 56.30 6516.077 FeII 2.891 -3.37 . . . 48.70 54.10 6154.226 NaI 2.102 -1.55 124.20 113.50 107.80 6160.747 NaI 2.104 -1.26 136.50 135.40 122.50 5711.088 MgI 4.346 -1.83 146.20 145.20 144.70 6318.717 MgI 5.108 -1.73 88.30 91.40 90.40 6319.237 MgI 5.108 -1.95 ...... 6319.495 MgI 5.108 -2.43 ...... 6696.023 AlI 3.140 -1.35 119.90 118.60 106.50 6698.673 AlI 3.140 -1.87 79.20 82.20 66.40 5690.425 SiI 4.929 -1.87 60.50 51.50 55.00 5701.104 SiI 4.929 -2.05 44.30 . . . 39.60 5772.146 SiI 5.082 -1.75 55.60 45.90 . . . 5793.073 SiI 4.929 -2.06 54.40 43.20 40.40 6142.483 SiI 5.619 -1.30 43.30 ...... 6145.016 SiI 5.616 -1.31 38.20 33.30 36.80 6155.134 SiI 5.619 -0.65 75.10 73.70 82.70 6243.815 SiI 5.616 -1.24 ...... 6244.466 SiI 5.616 -1.09 54.70 39.60 51.60 6161.297 CaI 2.523 -1.02 147.30 146.60 135.10 6166.439 CaI 2.521 -0.90 149.10 . . . 139.20 Continued on Next Page. . .

98 Table A.2 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 3336 6329 6329 6169.042 CaI 2.523 -0.55 ...... 6169.563 CaI 2.526 -0.27 ...... 6455.598 CaI 2.521 -1.32 136.50 139.10 119.80 6471.662 CaI 2.526 -0.59 ...... 6493.781 CaI 2.521 0.14 ...... 6499.650 CaI 2.523 -0.59 ...... 6572.779 CaI 0.000 -4.29 ...... 5657.895 ScII 1.507 -0.60 97.60 105.60 90.50 5684.190 ScII 1.507 -1.07 80.50 71.70 68.60 6245.621 ScII 1.507 -1.02 75.70 79.30 65.80 6300.684 ScII 1.507 -1.90 ...... 6320.850 ScII 1.500 -1.82 30.50 35.40 22.10 6604.599 ScII 1.357 -1.31 . . . 76.70 . . . 5702.656 TiI 2.291 -0.65 80.90 89.80 . . . 5716.445 TiI 2.297 -0.78 . . . 83.40 60.80 6064.626 TiI 1.046 -1.94 111.90 122.80 92.30 6091.171 TiI 2.267 -0.32 ...... 6092.792 TiI 1.887 -1.38 70.70 . . . 61.70 6261.099 TiI 1.429 -0.53 ...... 6303.757 TiI 1.443 -1.59 99.20 ...... 6312.236 TiI 1.460 -1.55 106.80 . . . 80.10 6336.099 TiI 1.443 -2.01 82.90 92.60 64.60 6599.105 TiI 0.900 -2.08 125.80 . . . 97.50 5418.768 TiII 1.581 -2.13 ...... 6491.561 TiII 2.061 -1.94 ...... 6559.588 TiII 2.048 -2.18 65.10 64.30 58.80 6606.949 TiII 2.061 -2.79 34.90 27.70 26.30 6680.135 TiII 3.095 -1.89 18.50 . . . 14.90 6039.713 VI 1.064 -0.65 120.90 129.90 103.30 6081.443 VI 1.051 -0.61 125.30 . . . 110.10 6090.194 VI 1.081 -0.07 145.60 . . . 125.70 6119.525 VI 1.064 -0.36 137.00 . . . 115.80 6135.365 VI 1.051 -0.76 ...... 102.80 6274.652 VI 0.267 -1.70 129.10 140.80 104.40 6285.160 VI 0.275 -1.54 134.00 155.90 110.40 6531.446 VI 1.218 -1.50 ...... 5760.828 NiI 4.105 -0.80 65.80 62.00 62.30 5805.213 NiI 4.167 -0.64 64.10 61.60 59.70 6176.812 NiI 4.088 -0.26 93.30 90.60 86.90 6204.604 NiI 4.088 -1.08 57.80 55.70 44.70 6223.984 NiI 4.105 -0.91 62.50 54.40 54.20 6378.260 NiI 4.153 -0.82 59.50 58.00 60.00 6127.475 ZrI 0.154 -1.06 73.40 93.50 50.60 6134.585 ZrI 0.000 -1.28 73.60 88.80 51.20 6140.535 ZrI 0.519 -1.41 ...... 6143.252 ZrI 0.071 -1.10 81.10 96.10 54.40 5853.668 BaII 0.604 -1.00 127.70 132.30 107.30 6141.713 BaII 0.704 -0.08 ...... 6496.897 BaII 0.604 -0.38 ...... 149.60

99 Table A.3. NGC 2420 Line list.

λ(A)˚ Element E.P. (eV) log(gf) 1418 4589 5050 5508 8015 8244 5417.037 FeI 4.415 -1.68 ...... 5441.339 FeI 4.312 -1.73 ...... 5445.042 FeI 4.386 -0.21 ...... 5466.396 FeI 4.371 -0.63 ...... 5466.987 FeI 3.573 -2.23 ...... 5505.881 FeI 4.412 -1.20 ...... 5522.446 FeI 4.209 -1.55 ...... 5557.897 FeI 3.111 -3.71 ...... 5576.089 FeI 3.430 -1.00 ...... 5679.024 FeI 4.652 -0.92 58.30 75.20 63.50 70.50 67.70 72.10 5731.762 FeI 4.256 -1.30 67.20 83.30 67.30 75.40 74.60 71.30 5752.032 FeI 4.548 -0.92 61.30 . . . 71.40 . . . 69.70 73.70 5760.344 FeI 3.642 -2.49 38.20 63.20 48.50 48.10 49.30 49.50 5775.081 FeI 4.220 -1.30 67.50 88.70 62.90 72.90 69.30 73.60 5778.453 FeI 2.588 -3.59 34.20 81.30 48.80 59.90 51.30 54.90 5793.915 FeI 4.220 -1.70 47.40 64.60 50.00 57.10 53.70 56.90 5806.726 FeI 4.607 -1.05 63.10 72.50 62.30 71.90 59.20 68.20 5855.076 FeI 4.607 -1.76 20.00 37.90 29.00 32.10 26.00 27.10 5856.088 FeI 4.294 -1.64 44.50 68.30 50.10 57.50 53.30 49.90 5916.247 FeI 2.458 -2.99 87.70 124.10 82.40 93.00 93.40 95.50 5927.789 FeI 4.652 -1.09 47.90 64.00 49.20 58.40 52.20 58.70 5929.677 FeI 4.548 -1.41 43.90 57.50 47.80 . . . 48.00 54.60 5934.655 FeI 3.928 -1.17 88.20 . . . 87.30 102.40 . . . 106.40 5987.065 FeI 4.795 -0.43 79.80 90.90 79.40 82.10 86.80 87.50 6027.051 FeI 4.076 -1.21 78.10 98.10 79.60 87.30 87.60 89.20 6056.004 FeI 4.733 -0.46 77.60 ...... 93.50 86.10 86.40 6079.008 FeI 4.652 -1.12 50.60 65.60 55.30 64.40 58.90 59.70 6093.643 FeI 4.607 -1.50 38.20 46.50 43.90 46.10 38.40 46.60 6096.664 FeI 3.984 -1.93 46.90 75.60 53.80 59.40 54.40 56.60 6151.617 FeI 2.176 -3.37 72.70 125.20 90.60 94.60 90.70 96.90 6165.360 FeI 4.143 -1.55 57.60 83.70 67.40 67.00 65.20 73.00 6180.203 FeI 2.727 -2.78 78.70 . . . 89.80 90.10 90.30 97.50 6187.989 FeI 3.943 -1.72 63.90 86.10 62.00 68.10 69.20 76.50 6200.313 FeI 2.608 -2.44 98.50 143.80 107.20 118.40 114.30 116.70 6229.226 FeI 2.845 -2.97 . . . 94.70 66.20 77.30 68.30 78.20 6232.641 FeI 3.654 -1.24 101.40 127.60 99.90 110.70 115.20 110.90 6240.645 FeI 2.223 -3.38 79.00 121.80 . . . 94.00 91.20 99.60 6246.318 FeI 3.602 -0.77 128.10 . . . 125.90 142.70 138.40 143.50 6270.223 FeI 2.858 -2.71 78.40 108.60 84.00 94.90 86.80 89.60 6301.501 FeI 3.654 -0.71 129.60 . . . 126.40 146.00 136.70 . . . 6336.824 FeI 3.686 -1.05 108.10 136.70 114.40 117.60 118.50 126.20 6393.600 FeI 2.433 -1.62 ...... 6475.624 FeI 2.559 -2.94 80.20 ...... 6481.869 FeI 2.279 -3.01 99.70 142.70 . . . 109.80 . . . 110.80 6494.980 FeI 2.404 -1.27 ...... 6498.940 FeI 0.958 -4.69 89.80 . . . 98.10 106.20 107.80 113.60 6533.928 FeI 4.558 -1.43 39.40 60.60 . . . 56.30 51.60 52.00 6546.238 FeI 2.758 -1.65 131.20 . . . 139.60 . . . 149.60 . . . 6569.214 FeI 4.733 -0.42 79.50 . . . 90.20 . . . 88.40 96.70 6575.016 FeI 2.588 -2.82 83.10 . . . 90.90 103.10 104.00 106.50 Continued on Next Page. . .

100 Table A.3 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 1418 4589 5050 5508 8015 8244 6592.913 FeI 2.727 -1.60 136.90 ...... 6593.870 FeI 2.433 -2.42 115.40 . . . 125.90 128.70 126.50 . . . 6597.559 FeI 4.795 -1.07 44.10 56.80 48.90 62.70 54.20 52.20 6608.025 FeI 2.279 -4.03 35.30 84.60 47.40 54.00 54.10 57.70 6609.110 FeI 2.559 -2.69 85.90 137.10 94.20 106.50 109.10 115.30 6609.678 FeI 0.990 -5.18 64.60 ...... 79.90 82.20 . . . 6627.545 FeI 4.548 -1.59 . . . 48.70 39.90 45.50 41.40 46.20 6646.931 FeI 2.608 -3.99 . . . 62.80 . . . 39.80 . . . 38.30 6677.985 FeI 2.692 -1.42 151.00 ...... 6703.566 FeI 2.758 -3.16 54.50 100.40 62.40 69.80 70.90 76.20 6705.102 FeI 4.607 -0.99 62.80 74.80 68.60 68.70 67.80 73.40 6710.318 FeI 1.485 -4.88 . . . 105.60 . . . 69.20 . . . 64.30 6713.743 FeI 4.795 -1.60 ...... 28.00 32.60 ...... 6725.357 FeI 4.103 -2.10 ...... 46.00 37.60 44.00 6726.666 FeI 4.607 -1.09 . . . 76.30 . . . 65.20 64.70 61.80 6733.150 FeI 4.638 -1.58 . . . 42.10 ...... 6793.258 FeI 4.076 -2.47 ...... 6810.262 FeI 4.607 -1.12 61.60 73.30 . . . 64.10 66.30 70.30 6820.371 FeI 4.638 -1.32 ...... 44.80 56.60 41.70 54.20 6828.592 FeI 4.638 -0.92 67.20 76.00 64.40 72.60 68.60 78.60 6839.830 FeI 2.559 -3.45 ...... 73.40 6842.685 FeI 4.638 -1.32 41.00 61.20 . . . 50.20 47.40 56.60 6843.655 FeI 4.548 -0.93 . . . 85.30 ...... 5425.249 FeII 3.199 -3.36 ...... 6084.103 FeII 3.199 -3.81 . . . 22.50 . . . 35.60 25.10 28.20 6149.246 FeII 3.889 -2.73 41.90 26.10 39.60 49.10 41.20 46.00 6247.559 FeII 3.891 -2.35 59.80 38.20 . . . 64.60 56.70 59.70 6369.459 FeII 2.891 -4.25 ...... 23.50 30.90 . . . 28.50 6416.919 FeII 3.891 -2.75 ...... 37.90 48.70 43.40 45.40 6432.677 FeII 2.891 -3.71 47.20 38.70 42.00 52.80 49.90 55.20 6456.381 FeII 3.903 -2.09 68.60 48.80 56.20 73.20 69.50 73.90 6516.077 FeII 2.891 -3.37 ...... 57.90 ...... 71.60 6154.226 NaI 2.102 -1.55 57.20 99.30 55.60 66.80 56.80 61.80 6160.747 NaI 2.104 -1.26 73.10 113.60 75.10 79.70 77.20 79.70 5711.088 MgI 4.346 -1.83 108.40 128.30 117.30 122.60 115.50 116.40 6318.717 MgI 5.108 -1.73 57.20 81.70 68.70 74.60 62.40 66.00 6319.237 MgI 5.108 -1.95 ...... 6319.495 MgI 5.108 -2.43 ...... 6696.023 AlI 3.140 -1.35 51.80 95.00 63.10 66.90 67.50 66.00 6698.673 AlI 3.140 -1.87 29.40 69.60 33.90 39.40 32.60 40.70 5690.425 SiI 4.929 -1.87 51.30 51.00 . . . 56.20 61.40 54.90 5701.104 SiI 4.929 -2.05 ...... 40.10 42.50 46.80 44.00 5772.146 SiI 5.082 -1.75 ...... 50.10 ...... 5793.073 SiI 4.929 -2.06 43.20 . . . 47.10 45.20 44.90 48.30 6142.483 SiI 5.619 -1.30 48.30 ...... 43.60 6145.016 SiI 5.616 -1.31 41.20 33.90 . . . 37.70 43.90 . . . 6155.134 SiI 5.619 -0.65 85.10 69.10 75.40 82.20 82.70 84.20 6243.815 SiI 5.616 -1.24 50.80 . . . 48.50 50.20 47.90 . . . 6244.466 SiI 5.616 -1.09 47.10 38.70 . . . 47.80 54.60 53.60 6161.297 CaI 2.523 -1.02 97.10 134.60 96.90 104.60 105.40 105.40 6166.439 CaI 2.521 -0.90 102.30 139.40 106.70 106.40 . . . 110.10 Continued on Next Page. . .

101 Table A.3 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 1418 4589 5050 5508 8015 8244 6169.042 CaI 2.523 -0.55 115.90 . . . 124.40 128.40 128.90 123.10 6169.563 CaI 2.526 -0.27 ...... 6455.598 CaI 2.521 -1.32 79.70 121.80 82.40 84.60 88.70 82.60 6471.662 CaI 2.526 -0.59 ...... 6493.781 CaI 2.521 0.14 ...... 6499.650 CaI 2.523 -0.59 124.70 . . . 118.90 132.40 129.90 124.20 6572.779 CaI 0.000 -4.29 76.90 ...... 98.30 . . . 5657.895 ScII 1.507 -0.60 . . . 98.60 80.10 . . . 95.60 97.80 5684.190 ScII 1.507 -1.07 65.40 ...... 60.30 ...... 6245.621 ScII 1.507 -1.02 63.50 70.50 61.00 58.50 63.50 73.20 6300.684 ScII 1.507 -1.90 ...... 17.20 ...... 6320.850 ScII 1.500 -1.82 ...... 27.90 29.40 26.20 6604.599 ScII 1.357 -1.31 56.80 . . . 62.20 . . . 66.70 69.10 5702.656 TiI 2.291 -0.65 ...... 34.40 . . . 44.70 33.00 5716.445 TiI 2.297 -0.78 . . . 77.60 21.70 35.40 . . . 26.20 6064.626 TiI 1.046 -1.94 ...... 47.20 56.00 48.70 49.30 6091.171 TiI 2.267 -0.32 39.30 . . . 57.80 55.30 55.40 53.00 6092.792 TiI 1.887 -1.38 ...... 24.30 . . . 29.80 25.30 6261.099 TiI 1.429 -0.53 86.80 . . . 96.10 104.60 109.30 107.30 6303.757 TiI 1.443 -1.59 . . . 111.30 . . . 40.30 42.40 31.90 6312.236 TiI 1.460 -1.55 ...... 36.30 38.70 . . . 36.20 6336.099 TiI 1.443 -2.01 . . . 84.30 ...... 24.80 19.30 6599.105 TiI 0.900 -2.08 ...... 54.00 49.30 53.40 51.10 5418.768 TiII 1.581 -2.13 ...... 6491.561 TiII 2.061 -1.94 58.80 85.40 59.00 72.80 72.50 65.80 6559.588 TiII 2.048 -2.18 45.90 65.80 45.80 56.60 55.00 54.60 6606.949 TiII 2.061 -2.79 . . . 36.30 ...... 26.20 22.00 6680.135 TiII 3.095 -1.89 . . . 19.80 . . . 20.40 ...... 6039.713 VI 1.064 -0.65 31.30 120.80 . . . 53.50 60.60 65.50 6081.443 VI 1.051 -0.61 37.90 . . . 61.90 52.20 61.80 59.40 6090.194 VI 1.081 -0.07 62.60 . . . 81.60 88.00 82.30 . . . 6119.525 VI 1.064 -0.36 54.50 142.20 . . . 75.50 72.90 79.50 6135.365 VI 1.051 -0.76 . . . 126.60 46.80 46.80 43.10 51.40 6274.652 VI 0.267 -1.70 . . . 133.70 50.40 ...... 50.80 6285.160 VI 0.275 -1.54 . . . 143.60 56.90 51.30 59.50 63.80 6531.446 VI 1.218 -1.50 ...... 5760.828 NiI 4.105 -0.80 ...... 48.80 52.20 49.40 53.00 5805.213 NiI 4.167 -0.64 40.70 53.70 46.90 58.60 45.50 55.00 6176.812 NiI 4.088 -0.26 66.30 84.20 75.30 79.30 81.70 81.80 6204.604 NiI 4.088 -1.08 . . . 39.10 35.60 39.20 38.00 35.10 6223.984 NiI 4.105 -0.91 38.20 50.60 44.00 48.10 . . . 45.60 6378.260 NiI 4.153 -0.82 42.10 52.20 37.10 52.40 44.40 49.10 6127.475 ZrI 0.154 -1.06 . . . 82.20 18.20 21.80 12.90 17.60 6134.585 ZrI 0.000 -1.28 8.00 78.40 14.40 23.00 . . . 16.70 6140.535 ZrI 0.519 -1.41 . . . 26.30 ...... 6143.252 ZrI 0.071 -1.10 7.90 88.40 21.00 23.30 13.20 22.70 5853.668 BaII 0.604 -1.00 103.70 147.90 99.20 117.80 119.70 112.90 6141.713 BaII 0.704 -0.08 ...... 6496.897 BaII 0.604 -0.38 ...... 128.70 ...... 152.30

102 Table A.4. NGC 6819 and NGC 6791 Line list.

λ(A)˚ Element E.P. (eV) log(gf) Cluster 0532 2111 2283 Cluster 7187 3426 5417.037 FeI 4.415 -1.68 N6819 ...... N6791 ...... 5441.339 FeI 4.312 -1.73 N6819 ...... N6791 ...... 5445.042 FeI 4.386 -0.21 N6819 ...... N6791 ...... 5466.396 FeI 4.371 -0.63 N6819 ...... N6791 ...... 5466.987 FeI 3.573 -2.23 N6819 ...... N6791 ...... 5505.881 FeI 4.412 -1.20 N6819 ...... N6791 ...... 5522.446 FeI 4.209 -1.55 N6819 ...... N6791 ...... 5557.897 FeI 3.111 -3.71 N6819 ...... N6791 ...... 5576.089 FeI 3.430 -1.00 N6819 ...... N6791 ...... 5679.024 FeI 4.652 -0.92 N6819 77.60 75.20 84.10 N6791 85.10 84.60 5731.762 FeI 4.256 -1.30 N6819 83.10 78.90 86.90 N6791 99.00 . . . 5752.032 FeI 4.548 -0.92 N6819 81.40 87.30 74.60 N6791 91.50 89.50 5760.344 FeI 3.642 -2.49 N6819 53.40 64.40 68.60 N6791 74.30 81.70 5775.081 FeI 4.220 -1.30 N6819 77.40 80.80 88.20 N6791 97.70 95.30 5778.453 FeI 2.588 -3.59 N6819 57.00 73.10 78.00 N6791 88.40 103.60 5793.915 FeI 4.220 -1.70 N6819 54.70 69.10 76.10 N6791 87.80 . . . 5806.726 FeI 4.607 -1.05 N6819 74.80 75.40 85.50 N6791 78.00 82.40 5855.076 FeI 4.607 -1.76 N6819 31.00 41.50 49.00 N6791 51.60 51.20 5856.088 FeI 4.294 -1.64 N6819 58.60 60.50 71.20 N6791 ...... 5916.247 FeI 2.458 -2.99 N6819 105.60 114.50 122.50 N6791 127.30 137.80 5927.789 FeI 4.652 -1.09 N6819 71.50 . . . 72.80 N6791 70.70 75.90 5929.677 FeI 4.548 -1.41 N6819 56.70 56.50 67.80 N6791 72.00 70.80 5934.655 FeI 3.928 -1.17 N6819 98.40 106.70 121.40 N6791 123.80 . . . 5987.065 FeI 4.795 -0.43 N6819 94.20 . . . 103.20 N6791 109.00 94.80 6027.051 FeI 4.076 -1.21 N6819 92.80 104.70 110.20 N6791 106.90 107.20 6056.004 FeI 4.733 -0.46 N6819 91.20 94.70 . . . N6791 ...... 6079.008 FeI 4.652 -1.12 N6819 63.10 69.10 79.90 N6791 74.10 73.80 6093.643 FeI 4.607 -1.50 N6819 44.50 60.10 60.10 N6791 . . . 64.20 6096.664 FeI 3.984 -1.93 N6819 60.90 64.80 78.10 N6791 74.60 79.30 6151.617 FeI 2.176 -3.37 N6819 95.70 109.80 115.60 N6791 134.60 145.40 6165.360 FeI 4.143 -1.55 N6819 70.10 81.30 81.40 N6791 99.70 89.60 6180.203 FeI 2.727 -2.78 N6819 100.10 101.30 114.30 N6791 137.50 136.10 6187.989 FeI 3.943 -1.72 N6819 68.60 83.70 88.10 N6791 103.40 . . . 6200.313 FeI 2.608 -2.44 N6819 115.40 134.40 142.60 N6791 ...... 6229.226 FeI 2.845 -2.97 N6819 74.50 82.50 94.70 N6791 . . . 118.00 6232.641 FeI 3.654 -1.24 N6819 . . . 130.00 131.70 N6791 131.80 138.60 6240.645 FeI 2.223 -3.38 N6819 102.40 104.60 117.50 N6791 128.20 140.50 6246.318 FeI 3.602 -0.77 N6819 149.10 155.20 . . . N6791 ...... 6270.223 FeI 2.858 -2.71 N6819 91.60 96.10 103.30 N6791 132.70 129.50 6301.501 FeI 3.654 -0.71 N6819 148.80 156.10 159.60 N6791 ...... 6336.824 FeI 3.686 -1.05 N6819 133.10 131.40 150.50 N6791 154.30 149.20 6393.600 FeI 2.433 -1.62 N6819 ...... N6791 ...... 6475.624 FeI 2.559 -2.94 N6819 ...... N6791 ...... 6481.869 FeI 2.279 -3.01 N6819 110.30 118.80 132.30 N6791 ...... 6494.980 FeI 2.404 -1.27 N6819 ...... N6791 ...... 6498.940 FeI 0.958 -4.69 N6819 113.20 120.80 136.90 N6791 159.50 . . . 6533.928 FeI 4.558 -1.43 N6819 ...... 69.50 N6791 67.10 64.70 6546.238 FeI 2.758 -1.65 N6819 ...... N6791 ...... 6569.214 FeI 4.733 -0.42 N6819 . . . 106.00 110.40 N6791 . . . 100.10 6575.016 FeI 2.588 -2.82 N6819 103.70 ...... N6791 ...... Continued on Next Page. . .

103 Table A.4 – Continued

λ(A)˚ Element E.P. (eV) log(gf) Cluster 0532 2111 2283 Cluster 7187 3426 6592.913 FeI 2.727 -1.60 N6819 ...... N6791 ...... 6593.870 FeI 2.433 -2.42 N6819 134.80 ...... N6791 ...... 6597.559 FeI 4.795 -1.07 N6819 60.90 63.20 69.80 N6791 61.40 64.60 6608.025 FeI 2.279 -4.03 N6819 56.10 73.40 80.50 N6791 97.40 . . . 6609.110 FeI 2.559 -2.69 N6819 110.90 124.60 130.30 N6791 149.60 . . . 6609.678 FeI 0.990 -5.18 N6819 85.40 104.10 . . . N6791 ...... 6627.545 FeI 4.548 -1.59 N6819 46.00 . . . 62.60 N6791 57.60 59.20 6646.931 FeI 2.608 -3.99 N6819 36.10 . . . 56.10 N6791 85.70 . . . 6677.985 FeI 2.692 -1.42 N6819 ...... N6791 ...... 6703.566 FeI 2.758 -3.16 N6819 77.50 83.70 92.40 N6791 96.60 113.80 6705.102 FeI 4.607 -0.99 N6819 70.00 75.90 83.60 N6791 79.30 . . . 6710.318 FeI 1.485 -4.88 N6819 63.20 81.50 89.30 N6791 117.00 . . . 6713.743 FeI 4.795 -1.60 N6819 36.10 ...... N6791 ...... 6725.357 FeI 4.103 -2.10 N6819 43.70 54.90 . . . N6791 56.80 56.80 6726.666 FeI 4.607 -1.09 N6819 67.10 ...... N6791 84.90 . . . 6733.150 FeI 4.638 -1.58 N6819 45.40 48.90 . . . N6791 66.30 58.60 6793.258 FeI 4.076 -2.47 N6819 27.50 40.20 . . . N6791 ...... 6810.262 FeI 4.607 -1.12 N6819 64.70 74.40 82.80 N6791 78.90 78.90 6820.371 FeI 4.638 -1.32 N6819 ...... N6791 ...... 6828.592 FeI 4.638 -0.92 N6819 73.10 86.40 84.00 N6791 98.30 87.40 6839.830 FeI 2.559 -3.45 N6819 80.50 87.40 98.50 N6791 116.50 116.50 6842.685 FeI 4.638 -1.32 N6819 63.50 56.00 67.50 N6791 74.20 67.30 6843.655 FeI 4.548 -0.93 N6819 83.90 90.00 . . . N6791 104.30 . . . 5425.249 FeII 3.199 -3.36 N6819 ...... N6791 ...... 6084.103 FeII 3.199 -3.81 N6819 29.60 25.60 32.40 N6791 20.40 16.40 6149.246 FeII 3.889 -2.73 N6819 ...... 39.10 N6791 25.20 . . . 6247.559 FeII 3.891 -2.35 N6819 56.50 59.70 53.60 N6791 . . . 27.40 6369.459 FeII 2.891 -4.25 N6819 . . . 22.80 . . . N6791 16.70 . . . 6416.919 FeII 3.891 -2.75 N6819 36.90 41.70 . . . N6791 ...... 6432.677 FeII 2.891 -3.71 N6819 48.20 44.80 46.00 N6791 30.10 28.20 6456.381 FeII 3.903 -2.09 N6819 61.20 64.20 63.60 N6791 45.80 34.60 6516.077 FeII 2.891 -3.37 N6819 ...... N6791 ...... 6154.226 NaI 2.102 -1.55 N6819 77.00 99.80 104.70 N6791 129.60 141.80 6160.747 NaI 2.104 -1.26 N6819 95.90 114.40 120.50 N6791 144.90 151.60 5711.088 MgI 4.346 -1.83 N6819 138.70 134.10 132.80 N6791 147.30 141.20 6318.717 MgI 5.108 -1.73 N6819 88.80 83.20 85.80 N6791 92.80 89.60 6319.237 MgI 5.108 -1.95 N6819 ...... N6791 ...... 6319.495 MgI 5.108 -2.43 N6819 ...... N6791 ...... 6696.023 AlI 3.140 -1.35 N6819 91.40 92.50 88.80 N6791 118.40 126.80 6698.673 AlI 3.140 -1.87 N6819 56.10 52.60 53.80 N6791 89.90 101.60 5690.425 SiI 4.929 -1.87 N6819 55.60 56.50 59.40 N6791 61.60 54.60 5701.104 SiI 4.929 -2.05 N6819 43.60 50.00 51.90 N6791 . . . 42.40 5772.146 SiI 5.082 -1.75 N6819 55.90 . . . 67.10 N6791 ...... 5793.073 SiI 4.929 -2.06 N6819 . . . 53.70 . . . N6791 ...... 6142.483 SiI 5.619 -1.30 N6819 45.60 ...... N6791 ...... 6145.016 SiI 5.616 -1.31 N6819 . . . 39.20 57.80 N6791 ...... 6155.134 SiI 5.619 -0.65 N6819 82.40 86.80 94.30 N6791 78.70 59.80 6243.815 SiI 5.616 -1.24 N6819 53.50 51.90 . . . N6791 ...... 6244.466 SiI 5.616 -1.09 N6819 . . . 61.00 65.50 N6791 52.40 44.50 6161.297 CaI 2.523 -1.02 N6819 115.70 . . . 130.40 N6791 ...... 6166.439 CaI 2.521 -0.90 N6819 115.10 132.40 133.30 N6791 137.50 . . . Continued on Next Page. . .

104 Table A.4 – Continued

λ(A)˚ Element E.P. (eV) log(gf) Cluster 0532 2111 2283 Cluster 7187 3426 6169.042 CaI 2.523 -0.55 N6819 . . . 144.50 . . . N6791 ...... 6169.563 CaI 2.526 -0.27 N6819 ...... N6791 ...... 6455.598 CaI 2.521 -1.32 N6819 103.30 112.60 112.40 N6791 . . . 155.80 6471.662 CaI 2.526 -0.59 N6819 140.20 146.90 . . . N6791 ...... 6493.781 CaI 2.521 0.14 N6819 ...... N6791 ...... 6499.650 CaI 2.523 -0.59 N6819 142.40 144.40 147.80 N6791 152.10 . . . 6572.779 CaI 0.000 -4.29 N6819 ...... N6791 ...... 5657.895 ScII 1.507 -0.60 N6819 91.10 94.10 95.40 N6791 101.70 65.40 5684.190 ScII 1.507 -1.07 N6819 62.30 60.50 . . . N6791 ...... 6245.621 ScII 1.507 -1.02 N6819 64.20 74.20 . . . N6791 90.40 . . . 6300.684 ScII 1.507 -1.90 N6819 21.40 19.90 . . . N6791 ...... 6320.850 ScII 1.500 -1.82 N6819 20.20 . . . 30.80 N6791 ...... 6604.599 ScII 1.357 -1.31 N6819 61.20 62.20 71.20 N6791 . . . 47.30 5702.656 TiI 2.291 -0.65 N6819 44.50 . . . 70.20 N6791 99.50 116.40 5716.445 TiI 2.297 -0.78 N6819 31.00 44.40 60.80 N6791 94.90 109.70 6064.626 TiI 1.046 -1.94 N6819 43.30 76.50 89.10 N6791 ...... 6091.171 TiI 2.267 -0.32 N6819 60.20 72.70 . . . N6791 116.70 . . . 6092.792 TiI 1.887 -1.38 N6819 25.60 45.60 49.30 N6791 93.20 121.50 6261.099 TiI 1.429 -0.53 N6819 110.30 ...... N6791 ...... 6303.757 TiI 1.443 -1.59 N6819 42.20 . . . 73.70 N6791 ...... 6312.236 TiI 1.460 -1.55 N6819 50.40 67.00 . . . N6791 ...... 6336.099 TiI 1.443 -2.01 N6819 . . . 45.00 56.10 N6791 ...... 6599.105 TiI 0.900 -2.08 N6819 49.70 73.50 . . . N6791 ...... 5418.768 TiII 1.581 -2.13 N6819 ...... N6791 ...... 6491.561 TiII 2.061 -1.94 N6819 59.50 67.50 72.70 N6791 . . . 81.40 6559.588 TiII 2.048 -2.18 N6819 47.90 . . . 65.90 N6791 64.60 . . . 6606.949 TiII 2.061 -2.79 N6819 20.40 25.10 28.10 N6791 30.20 . . . 6680.135 TiII 3.095 -1.89 N6819 ...... N6791 ...... 6039.713 VI 1.064 -0.65 N6819 63.20 . . . 98.10 N6791 143.90 158.40 6081.443 VI 1.051 -0.61 N6819 75.50 87.20 . . . N6791 . . . 156.90 6090.194 VI 1.081 -0.07 N6819 94.20 111.30 125.90 N6791 ...... 6119.525 VI 1.064 -0.36 N6819 88.90 103.70 114.90 N6791 ...... 6135.365 VI 1.051 -0.76 N6819 62.50 76.60 94.20 N6791 ...... 6274.652 VI 0.267 -1.70 N6819 . . . 91.30 97.30 N6791 146.90 . . . 6285.160 VI 0.275 -1.54 N6819 72.50 102.70 107.70 N6791 ...... 6531.446 VI 1.218 -1.50 N6819 ...... N6791 ...... 5760.828 NiI 4.105 -0.80 N6819 53.50 54.50 65.50 N6791 67.40 63.70 5805.213 NiI 4.167 -0.64 N6819 57.50 57.90 62.60 N6791 66.80 65.90 6176.812 NiI 4.088 -0.26 N6819 79.10 91.80 91.60 N6791 91.60 91.40 6204.604 NiI 4.088 -1.08 N6819 40.90 50.70 49.30 N6791 55.10 53.30 6223.984 NiI 4.105 -0.91 N6819 . . . 53.40 54.60 N6791 . . . 57.60 6378.260 NiI 4.153 -0.82 N6819 55.30 55.60 61.40 N6791 . . . 60.20 6127.475 ZrI 0.154 -1.06 N6819 21.90 35.80 39.40 N6791 103.40 129.70 6134.585 ZrI 0.000 -1.28 N6819 25.40 33.40 35.40 N6791 ...... 6140.535 ZrI 0.519 -1.41 N6819 ...... N6791 . . . 92.60 6143.252 ZrI 0.071 -1.10 N6819 24.90 44.80 42.80 N6791 112.60 134.30 5853.668 BaII 0.604 -1.00 N6819 101.40 112.80 121.40 N6791 131.70 146.70 6141.713 BaII 0.704 -0.08 N6819 ...... N6791 ...... 6496.897 BaII 0.604 -0.38 N6819 139.00 ...... N6791 ......

105 Table A.5. NGC 7789 Line list.

λ(A)˚ Element E.P. (eV) log(gf) 0125 0586 1221 1459 7208 5417.037 FeI 4.415 -1.68 ...... 5441.339 FeI 4.312 -1.73 66.20 60.40 67.60 . . . 56.10 5445.042 FeI 4.386 -0.21 135.50 134.40 145.80 . . . 126.90 5466.396 FeI 4.371 -0.63 115.40 110.50 121.40 108.30 107.20 5466.987 FeI 3.573 -2.23 83.30 77.30 91.20 81.40 74.60 5505.881 FeI 4.412 -1.20 ...... 5522.446 FeI 4.209 -1.55 76.70 73.90 83.10 ...... 5557.897 FeI 3.111 -3.71 ...... 47.40 ...... 5576.089 FeI 3.430 -1.00 ...... 151.70 141.90 5679.024 FeI 4.652 -0.92 78.50 80.40 87.40 82.80 75.20 5731.762 FeI 4.256 -1.30 90.70 82.90 100.30 84.80 79.60 5752.032 FeI 4.548 -0.92 87.40 85.70 95.90 83.20 . . . 5760.344 FeI 3.642 -2.49 68.40 60.10 79.60 72.70 59.00 5775.081 FeI 4.220 -1.30 86.90 86.80 95.90 88.80 84.30 5778.453 FeI 2.588 -3.59 85.20 80.20 90.70 79.60 66.50 5793.915 FeI 4.220 -1.70 67.50 64.00 76.50 70.00 67.30 5806.726 FeI 4.607 -1.05 83.20 76.70 92.70 77.70 79.10 5855.076 FeI 4.607 -1.76 45.90 40.00 50.90 48.40 39.50 5856.088 FeI 4.294 -1.64 66.60 67.30 74.10 65.90 61.90 5916.247 FeI 2.458 -2.99 122.00 111.60 130.80 117.00 108.60 5927.789 FeI 4.652 -1.09 70.00 72.20 86.40 70.20 69.00 5929.677 FeI 4.548 -1.41 66.70 67.00 68.80 62.60 59.70 5934.655 FeI 3.928 -1.17 111.90 116.90 122.40 112.40 104.50 5987.065 FeI 4.795 -0.43 . . . 95.80 105.10 105.40 95.70 6027.051 FeI 4.076 -1.21 101.90 ...... 6056.004 FeI 4.733 -0.46 97.90 97.60 111.70 97.40 94.20 6079.008 FeI 4.652 -1.12 72.60 71.40 81.80 68.80 68.30 6093.643 FeI 4.607 -1.50 60.30 56.90 64.60 51.80 57.20 6096.664 FeI 3.984 -1.93 72.80 74.90 79.20 71.30 70.10 6151.617 FeI 2.176 -3.37 122.60 118.00 131.40 114.90 104.30 6165.360 FeI 4.143 -1.55 84.80 81.30 98.60 83.20 72.70 6180.203 FeI 2.727 -2.78 ...... 113.80 106.30 6187.989 FeI 3.943 -1.72 86.30 89.00 97.40 82.70 81.40 6200.313 FeI 2.608 -2.44 143.60 135.10 . . . 134.30 122.80 6229.226 FeI 2.845 -2.97 102.80 91.40 104.70 91.70 90.50 6232.641 FeI 3.654 -1.24 138.40 123.60 144.60 125.80 121.50 6240.645 FeI 2.223 -3.38 116.10 114.80 133.00 114.50 107.40 6246.318 FeI 3.602 -0.77 ...... 6270.223 FeI 2.858 -2.71 ...... 121.80 106.20 101.40 6301.501 FeI 3.654 -0.71 ...... 6336.824 FeI 3.686 -1.05 142.20 140.80 152.10 140.60 133.60 6393.600 FeI 2.433 -1.62 ...... 6475.624 FeI 2.559 -2.94 125.50 111.80 137.90 119.30 112.50 6481.869 FeI 2.279 -3.01 135.00 133.80 . . . 128.50 123.10 6494.980 FeI 2.404 -1.27 ...... 6498.940 FeI 0.958 -4.69 ...... 130.90 123.30 6533.928 FeI 4.558 -1.43 66.50 59.50 76.90 60.20 59.80 6546.238 FeI 2.758 -1.65 ...... 6569.214 FeI 4.733 -0.42 111.20 ...... 104.50 101.00 6575.016 FeI 2.588 -2.82 ...... 116.30 Continued on Next Page. . .

106 Table A.5 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 0125 0586 1221 1459 7208 6592.913 FeI 2.727 -1.60 ...... 6593.870 FeI 2.433 -2.42 . . . 152.60 ...... 144.90 6597.559 FeI 4.795 -1.07 66.90 67.10 74.60 63.70 60.90 6608.025 FeI 2.279 -4.03 87.40 77.10 90.80 . . . 68.30 6609.110 FeI 2.559 -2.69 138.20 131.60 . . . 130.00 124.40 6609.678 FeI 0.990 -5.18 118.00 108.40 130.10 108.50 90.90 6627.545 FeI 4.548 -1.59 ...... 6646.931 FeI 2.608 -3.99 ...... 6677.985 FeI 2.692 -1.42 ...... 6703.566 FeI 2.758 -3.16 ...... 6705.102 FeI 4.607 -0.99 ...... 6710.318 FeI 1.485 -4.88 ...... 6713.743 FeI 4.795 -1.60 ...... 6725.357 FeI 4.103 -2.10 ...... 6726.666 FeI 4.607 -1.09 ...... 6733.150 FeI 4.638 -1.58 ...... 6793.258 FeI 4.076 -2.47 ...... 6810.262 FeI 4.607 -1.12 ...... 6820.371 FeI 4.638 -1.32 ...... 6828.592 FeI 4.638 -0.92 ...... 6839.830 FeI 2.559 -3.45 ...... 6842.685 FeI 4.638 -1.32 ...... 6843.655 FeI 4.548 -0.93 ...... 5425.249 FeII 3.199 -3.36 43.60 45.10 . . . 45.90 52.40 6084.103 FeII 3.199 -3.81 23.30 26.40 28.40 22.00 30.80 6149.246 FeII 3.889 -2.73 30.20 34.30 35.10 33.80 42.20 6247.559 FeII 3.891 -2.35 49.10 48.90 55.10 49.80 56.40 6369.459 FeII 2.891 -4.25 22.90 . . . 24.60 22.80 29.10 6416.919 FeII 3.891 -2.75 ...... 44.00 6432.677 FeII 2.891 -3.71 43.40 41.50 50.90 39.70 51.20 6456.381 FeII 3.903 -2.09 54.30 64.00 62.20 60.70 69.10 6516.077 FeII 2.891 -3.37 ...... 60.30 . . . 6154.226 NaI 2.102 -1.55 97.30 94.10 104.30 88.30 82.70 6160.747 NaI 2.104 -1.26 118.60 116.60 116.90 111.30 95.00 5711.088 MgI 4.346 -1.83 132.10 142.70 135.10 132.10 126.30 6318.717 MgI 5.108 -1.73 85.10 95.90 85.20 77.60 72.60 6319.237 MgI 5.108 -1.95 ...... 6319.495 MgI 5.108 -2.43 ...... 6696.023 AlI 3.140 -1.35 ...... 6698.673 AlI 3.140 -1.87 ...... 5690.425 SiI 4.929 -1.87 59.80 59.20 62.90 62.10 61.00 5701.104 SiI 4.929 -2.05 51.30 56.30 53.20 47.70 54.00 5772.146 SiI 5.082 -1.75 . . . 61.70 . . . 59.60 60.50 5793.073 SiI 4.929 -2.06 52.40 53.40 . . . 55.20 . . . 6142.483 SiI 5.619 -1.30 41.60 . . . 47.30 53.80 . . . 6145.016 SiI 5.616 -1.31 . . . 47.90 44.90 53.40 44.00 6155.134 SiI 5.619 -0.65 81.20 84.40 86.80 81.70 87.60 6243.815 SiI 5.616 -1.24 ...... 6244.466 SiI 5.616 -1.09 62.20 56.80 52.90 60.90 64.10 6161.297 CaI 2.523 -1.02 131.00 128.60 139.50 129.50 115.80 6166.439 CaI 2.521 -0.90 128.80 125.80 138.90 128.80 119.90 Continued on Next Page. . .

107 Table A.5 – Continued

λ(A)˚ Element E.P. (eV) log(gf) 0125 0586 1221 1459 7208 6169.042 CaI 2.523 -0.55 . . . 150.10 ...... 135.90 6169.563 CaI 2.526 -0.27 ...... 6455.598 CaI 2.521 -1.32 115.10 112.90 122.10 116.70 106.80 6471.662 CaI 2.526 -0.59 . . . 147.00 . . . 147.40 144.50 6493.781 CaI 2.521 0.14 ...... 6499.650 CaI 2.523 -0.59 150.10 142.80 ...... 134.40 6572.779 CaI 0.000 -4.29 ...... 5657.895 ScII 1.507 -0.60 101.00 90.80 110.40 102.60 93.30 5684.190 ScII 1.507 -1.07 77.20 69.30 ...... 61.10 6245.621 ScII 1.507 -1.02 . . . 72.80 90.60 ...... 6300.684 ScII 1.507 -1.90 ...... 40.90 22.30 6320.850 ScII 1.500 -1.82 . . . 35.50 39.10 42.50 27.10 6604.599 ScII 1.357 -1.31 78.70 . . . 92.00 80.80 . . . 5702.656 TiI 2.291 -0.65 66.90 61.50 ...... 36.60 5716.445 TiI 2.297 -0.78 61.40 52.20 . . . 53.60 . . . 6064.626 TiI 1.046 -1.94 102.60 87.00 106.80 80.30 63.20 6091.171 TiI 2.267 -0.32 92.30 80.70 . . . 87.70 65.40 6092.792 TiI 1.887 -1.38 58.70 52.40 66.10 51.00 33.60 6261.099 TiI 1.429 -0.53 ...... 6303.757 TiI 1.443 -1.59 ...... 102.50 81.10 . . . 6312.236 TiI 1.460 -1.55 ...... 53.60 6336.099 TiI 1.443 -2.01 ...... 6599.105 TiI 0.900 -2.08 108.40 88.70 111.20 90.20 69.30 5418.768 TiII 1.581 -2.13 87.30 84.20 . . . 84.40 . . . 6491.561 TiII 2.061 -1.94 . . . 73.20 80.80 . . . 64.50 6559.588 TiII 2.048 -2.18 57.60 49.70 63.20 54.90 49.10 6606.949 TiII 2.061 -2.79 . . . 22.60 ...... 20.70 6680.135 TiII 3.095 -1.89 ...... 6039.713 VI 1.064 -0.65 . . . 92.80 . . . 102.20 . . . 6081.443 VI 1.051 -0.61 113.60 105.90 127.40 97.20 82.20 6090.194 VI 1.081 -0.07 134.30 . . . 145.50 . . . 104.80 6119.525 VI 1.064 -0.36 129.10 113.20 137.90 118.10 93.20 6135.365 VI 1.051 -0.76 113.00 97.70 109.90 97.70 78.90 6274.652 VI 0.267 -1.70 121.20 101.00 124.80 96.50 75.40 6285.160 VI 0.275 -1.54 129.00 111.30 142.30 115.20 96.80 6531.446 VI 1.218 -1.50 ...... 5760.828 NiI 4.105 -0.80 64.20 64.80 72.40 59.20 55.80 5805.213 NiI 4.167 -0.64 62.80 62.10 69.70 63.10 56.20 6176.812 NiI 4.088 -0.26 86.80 87.30 101.20 89.90 88.80 6204.604 NiI 4.088 -1.08 48.10 48.60 56.90 52.20 47.30 6223.984 NiI 4.105 -0.91 54.20 53.10 63.20 49.80 55.70 6378.260 NiI 4.153 -0.82 60.80 57.20 66.80 55.60 56.60 6127.475 ZrI 0.154 -1.06 60.80 49.80 75.60 44.30 32.90 6134.585 ZrI 0.000 -1.28 61.40 45.60 70.30 48.70 33.00 6140.535 ZrI 0.519 -1.41 20.90 . . . 22.10 ...... 6143.252 ZrI 0.071 -1.10 71.10 59.60 75.00 54.30 32.30 5853.668 BaII 0.604 -1.00 147.70 127.70 148.90 133.20 128.10 6141.713 BaII 0.704 -0.08 ...... 6496.897 BaII 0.604 -0.38 ......

108 Appendix B

Abundance Measurements Derived for 20 Open Clusters

Here we present the detailed per star abundance determinations for the rest of the non- calibration star sample used in this study.

109 Table B.1. Measured Abundances for ASCC 14 and Collinder 106.

ASCC 14 Col 106 Col 106

Ratio 0310 2571 7411

[FeI/H] 0.15±0.06 −0.17±0.07 −0.34±0.07 [FeII/H] 0.14±0.06 −0.13±0.04 −0.36±0.07 [Na/Fe] 0.24±0.06 −0.20±0.02 0.08±0.03 [Mg/Fe] −0.08±0.03 −0.03±0.07 −0.06±0.08 [Al/Fe] ...... [Si/Fe] 0.07±0.08 0.20±0.05 0.15±0.05 [Ca/Fe] −0.40±0.05 −0.24±0.03 −0.33±0.06 [ScII/Fe] −0.07±0.02 0.02±0.06 0.03±0.03 [Ti/Fe] −0.04±0.08 0.11±0.02 0.21±0.07 [TiII/Fe] 0.01±0.03 0.10±0.08 0.23±0.07 [V/Fe] −0.12±0.08 −0.02±0.08 −0.12±0.07 [Ni/Fe] 0.00±0.09 0.00±0.04 0.02±0.08 [Zr/Fe] 0.05±0.04 −0.04±0.05 −0.17±0.07 [BaII/Fe] 0.16±0.01 0.00±0.00 0.22±0.01 [YII/Fe] 0.17±0.01 0.20±0.01 0.32±0.02 [ZrII/Fe] 0.07±0.00 . . . 0.15±0.00 [LaII/Fe] 0.05±0.02 0.18±0.01 0.15±0.02 [CeII/Fe] 0.13±0.01 0.19±0.01 0.47±0.01 [NdII/Fe] 0.12±0.01 0.31±0.01 0.44±0.01 [EuII/Fe] −0.05±0.01 0.18±0.02 0.12±0.02

Table B.2. Measured Abundances for NGC 1912, Berkeley 52, and FSR 489.

NGC 1912 Be 53 FSR 498

Ratio 5224 3206 5307

[FeI/H] −0.15±0.08 −0.13±0.05 −0.30±0.01 [FeII/H] −0.15±0.08 −0.14±0.04 −0.30±0.07 [Na/Fe] 0.22±0.01 −0.18±0.08 −0.08±0.00 [Mg/Fe] −0.01±0.09 −0.12±0.06 0.10±0.07 [Al/Fe] ...... [Si/Fe] 0.19±0.09 0.12±0.02 0.18±0.07 [Ca/Fe] −0.20±0.08 −0.09±0.06 −0.15±0.04 [ScII/Fe] . . . 0.23±0.08 0.10±0.04 [Ti/Fe] 0.01±0.06 0.07±0.06 0.02±0.02 [TiII/Fe] 0.02±0.08 0.16±0.01 0.10±0.01 [Ni/Fe] −0.04±0.02 −0.02±0.03 −0.01±0.05 [Zr/Fe] 0.37±0.00 −0.37±0.00 −0.02±0.07 [BaII/Fe] 0.48±0.00 0.58±0.00 0.36±0.00 [LaII/Fe] 0.18±0.02 0.28±0.02 0.08±0.01 [CeII/Fe] 0.08±0.02 0.10±0.02 0.15±0.01 [EuII/Fe] −0.15±0.02 0.21±0.01 0.00±0.02

110 Table B.3. Measured Abundances for NGC 6705.

Ratio 0414 2435 4564 7119

[FeI/H] 0.05±0.07 0.03±0.07 0.02±0.06 0.07±0.05 [FeII/H] 0.05±0.07 0.03±0.06 0.00±0.07 0.07±0.03 [Na/Fe] 0.47±0.00 0.38±0.04 0.28±0.06 0.30±0.03 [Mg/Fe] −0.12±0.01 0.01±0.00 −0.46±0.05 −0.05±0.02 [Al/Fe] 0.17±0.07 0.06±0.06 −0.01±0.02 0.14±0.08 [Si/Fe] −0.01±0.05 0.18±0.03 −0.02±0.05 0.02±0.06 [Ca/Fe] −0.08±0.01 −0.15±0.04 0.03±0.06 −0.10±0.08 [ScII/Fe] −0.18±0.04 0.10±0.02 −0.20±0.06 0.04±0.06 [Ti/Fe] −0.10±0.07 −0.03±0.09 −0.07±0.05 −0.06±0.08 [TiII/Fe] −0.11±0.01 −0.03±0.03 −0.06±0.05 0.03±0.01 [V/Fe] −0.02±0.03 0.01±0.08 −0.09±0.05 −0.07±0.04 [Ni/Fe] −0.04±0.06 −0.03±0.05 −0.03±0.07 0.03±0.08 [Zr/Fe] −0.10±0.06 −0.02±0.02 −0.13±0.08 −0.04±0.06 [BaII/Fe] 0.14±0.00 0.22±0.00 0.42±0.00 0.62±0.00 [YII/Fe] −0.12±0.01 −0.01±0.02 −0.07±0.02 0.02±0.01 [ZrII/Fe] ...... −0.14±0.01 0.03±0.02 [LaII/Fe] 0.10±0.01 0.13±0.01 0.01±0.01 0.14±0.02 [CeII/Fe] 0.21±0.01 0.06±0.01 0.11±0.02 0.21±0.01 [NdII/Fe] 0.13±0.01 0.09±0.01 0.13±0.02 0.16±0.01 [EuII/Fe] 0.03±0.02 0.11±0.01 −0.11±0.02 0.08±0.01

Table B.4. Measured Abundances for NGC 1896.

Ratio 0425 1527 6170

[FeI/H] −0.09±0.07 −0.07±0.07 −0.12±0.01 [FeII/H] −0.08±0.07 −0.05±0.06 −0.11±0.09 [Na/Fe] 0.04±0.03 0.07±0.02 0.15±0.05 [Mg/Fe] −0.02±0.07 −0.03±0.06 0.02±0.08 [Al/Fe] −0.07±0.06 −0.08±0.07 0.08±0.08 [Si/Fe] 0.01±0.08 −0.02±0.07 0.05±0.05 [Ca/Fe] −0.05±0.08 −0.03±0.05 −0.01±0.05 [ScII/Fe] −0.08±0.08 −0.05±0.07 −0.01±0.05 [Ti/Fe] 0.05±0.07 0.02±0.08 −0.01±0.09 [TiII/Fe] 0.06±0.07 0.03±0.08 0.00±0.09 [V/Fe] −0.03±0.08 0.01±0.09 −0.07±0.08 [Ni/Fe] −0.01±0.05 −0.04±0.06 0.01±0.07 [Zr/Fe] −0.06±0.05 0.05±0.05 −0.01±0.07 [BaII/Fe] 0.45±0.00 0.48±0.00 0.36±0.00 [LaII/Fe] 0.05±0.02 0.09±0.02 0.13±0.02 [EuII/Fe] −0.01±0.02 −0.05±0.02 0.01±0.02

111 Table B.5. Measured Abundances for IC 1369.

Ratio 4158 4571 5175 5256

[FeI/H] −0.030.05 −0.030.05 0.050.06 0.020.06 [FeII/H] −0.020.05 0.000.06 0.000.05 0.040.04 [Na/Fe] 0.260.04 0.200.03 0.150.02 0.150.04 [Mg/Fe] −0.070.06 −0.100.06 0.010.00 0.000.00 [Al/Fe] . . . −0.100.05 ...... [Si/Fe] −0.040.07 0.020.05 −0.010.06 0.050.04 [Ca/Fe] 0.020.10 0.000.08 −0.190.09 −0.210.07 [ScII/Fe] 0.020.05 0.050.05 0.000.08 0.020.08 [Ti/Fe] 0.020.08 −0.020.02 0.000.08 0.000.08 [TiII/Fe] 0.070.08 0.030.02 −0.040.09 0.010.05 [VI/Fe] −0.010.08 −0.020.07 0.060.07 −0.030.10 [Ni/Fe] −0.060.08 0.000.07 −0.030.08 −0.060.05 [Zr/Fe] 0.050.07 0.060.02 0.050.09 −0.130.09 [BaII/Fe] 0.590.00 0.630.00 0.260.00 0.560.00 [LaII/Fe] 0.050.01 0.190.02 0.200.01 0.230.02 [EuII/Fe] −0.020.02 −0.010.02 −0.040.02 −0.010.02

Table B.6. Measured Abundances for King 5 and NGC 2355.

King 5 King 5 NGC 2355

Ratio 7355 7511 1307

[FeI/H] −0.13±0.08 −0.15±0.07 0.13±0.05 [FeII/H] −0.14±0.06 −0.16±0.04 0.13±0.07 [Na/Fe] 0.09±0.06 0.02±0.06 0.31±0.03 [Mg/Fe] −0.20±0.02 −0.14±0.05 0.00±0.00 [Al/Fe] ...... 0.13±0.02 [Si/Fe] 0.20±0.08 0.13±0.07 0.04±0.08 [Ca/Fe] −0.22±0.04 −0.18±0.07 −0.25±0.05 [ScII/Fe] 0.07±0.08 0.03±0.07 −0.08±0.05 [Ti/Fe] −0.03±0.08 0.15±0.05 −0.02±0.04 [TiII/Fe] −0.03±0.04 0.14±0.05 0.00±0.05 [V/Fe] −0.07±0.04 −0.01±0.08 0.15±0.06 [Ni/Fe] −0.03±0.07 −0.04±0.07 −0.02±0.06 [Zr/Fe] −0.07±0.07 0.15±0.05 −0.02±0.09 [BaII/Fe] 0.25±0.00 0.66±0.00 0.42±0.00 [LaII/Fe] 0.22±0.01 0.30±0.02 0.16±0.01 [EuII/Fe] 0.19±0.01 0.10±0.02 −0.14±0.01

112 Table B.7. Measured Abundances for Berkeley 9, Berkeley 19, and Berkeley 31.

Be 9 Be 19 Be 19 Be 31

Ratio 9538 7217 4330 6069

[FeI/H] −0.07±0.06 −0.17±0.08 −0.26±0.07 −0.27±0.04 [FeII/H] −0.07±0.09 −0.13±0.04 −0.23±0.06 −0.25±0.04 [Na/Fe] 0.00±0.06 0.09±0.03 0.08±0.08 0.11±0.08 [Mg/Fe] −0.22±0.09 0.00±0.02 −0.01±0.07 −0.15±0.06 [Al/Fe] 0.23±0.09 ...... 0.07±0.47 [Si/Fe] 0.06±0.03 0.17±0.07 0.14±0.05 0.12±0.05 [Ca/Fe] −0.36±0.08 −0.22±0.02 −0.14±0.05 −0.17±0.15 [ScII/Fe] 0.04±0.08 0.04±0.05 0.05±0.05 −0.08±0.09 [Ti/Fe] 0.04±0.09 0.06±0.08 0.01±0.03 0.08±0.09 [TiII/Fe] 0.09±0.02 0.02±0.09 0.03±0.09 0.07±0.08 [V/Fe] 0.08±0.05 0.03±0.06 0.04±0.07 0.13±0.07 [Ni/Fe] −0.03±0.03 −0.03±0.08 −0.03±0.04 0.05±0.05 [Zr/Fe] 0.25±0.00 −0.01±0.06 −0.06±0.04 −0.07±0.06 [BaII/Fe] 0.18±0.08 0.51±0.00 0.31±0.00 −0.07±0.00 [LaII/Fe] 0.14±0.02 0.13±0.02 0.16±0.02 0.13±0.02 [EuII/Fe] −0.07±0.01 −0.17±0.01 0.06±0.02 0.11±0.01

Table B.8. Measured Abundances for Melotte 71 and NGC 1798.

Mel 71 Mel 71 NGC 1798 NGC 1798

Ratio 9538 7217 4330 6069

[FeI/H] −0.25±0.07 −0.28±0.07 −0.18±0.08 −0.14±0.09 [FeII/H] −0.26±0.07 −0.28±0.07 −0.20±0.08 −0.15±0.09 [Na/Fe] 0.15±0.08 0.00±0.08 0.13±0.07 −0.00±0.08 [Mg/Fe] 0.02±0.04 0.00±0.05 −0.05±0.07 −0.08±0.07 [Al/Fe] 0.13±0.06 0.10±0.06 −0.01±0.08 0.04±0.08 [Si/Fe] 0.09±0.06 0.13±0.06 0.01±0.06 0.02±0.06 [Ca/Fe] −0.13±0.05 −0.10±0.05 −0.01±0.10 −0.03±0.09 [ScII/Fe] 0.07±0.08 0.09±0.08 −0.13±0.09 −0.08±0.08 [Ti/Fe] −0.15±0.09 −0.13±0.09 −0.05±0.08 −0.03±0.08 [TiII/Fe] −0.13±0.05 −0.10±0.06 −0.06±0.08 −0.03±0.07 [V/Fe] −0.11±0.06 −0.15±0.05 −0.10±0.07 0.08±0.07 [Ni/Fe] −0.04±0.06 0.00±0.07 −0.00±0.06 0.01±0.06 [Zr/Fe] 0.08±0.07 −0.01±0.08 0.22±0.04 −0.02±0.05 [BaII/Fe] 0.20±0.00 0.21±0.00 0.32±0.00 0.35±0.00 [LaII/Fe] 0.20±0.01 0.23±0.01 0.09±0.01 0.10±0.02 [EuII/Fe] 0.17±0.02 0.21±0.01 0.02±0.02 0.04±0.01

113 Table B.9. Measured Abundances for NGC 6811 and NGC 7062.

NGC 6811 NGC 6811 NGC 7062

Ratio 2426 4098 2136

[FeI/H] 0.04±0.06 0.00±0.06 0.02±0.08 [FeII/H] 0.04±0.04 0.00±0.02 0.03±0.07 [Na/Fe] 0.04±0.02 0.09±0.04 0.15±0.04 [Mg/Fe] 0.05±0.03 −0.09±0.04 −0.05±0.05 [Al/Fe] ...... [Si/Fe] −0.02±0.07 0.00±0.00 0.00±0.07 [Ca/Fe] −0.15±0.08 0.01±0.04 −0.03±0.09 [ScII/Fe] ...... −0.14±0.06 [Ti/Fe] −0.10±0.08 −0.06±0.02 −0.07±0.07 [TiII/Fe] −0.10±0.06 −0.09±0.08 −0.04±0.08 [V/Fe] ...... −0.07±0.07 [Ni/Fe] −0.09±0.04 −0.03±0.07 −0.05±0.06 [Zr/Fe] 0.12±0.02 0.16±0.03 0.19±0.05 [BaII/Fe] 0.54±0.00 0.38±0.00 0.38±0.01 [LaII/Fe] 0.06±0.01 0.26±0.02 0.06±0.01 [EuII/Fe] −0.01±0.01 −0.06±0.01 −0.08±0.01

Table B.10. Measured Abundances for Ruprecht 24.

Ratio 4132 7237

[FeI/H] −0.40±0.09 −0.32±0.09 [FeII/H] −0.38±0.04 −0.30±0.09 [Na/Fe] 0.09±0.01 0.05±0.01 [Mg/Fe] 0.10±0.06 0.03±0.09 [Al/Fe] ...... [Si/Fe] 0.18±0.05 0.06±0.00 [Ca/Fe] −0.13±0.08 −0.08±0.08 [ScII/Fe] 0.16±0.05 0.26±0.05 [Ti/Fe] 0.07±0.04 0.13±0.06 [TiII/Fe] 0.15±0.04 0.16±0.06 [V/Fe] 0.01±0.07 −0.07±0.08 [Ni/Fe] 0.05±0.05 0.01±0.09 [Zr/Fe] 0.12±0.09 −0.02±0.09 [BaII/Fe] 0.21±0.01 0.39±0.00 [LaII/Fe] 0.17±0.01 0.16±0.01 [EuII/Fe] 0.33±0.01 0.37±0.01

114 Table B.11. Measured Abundances for Berkeley 17.

Ratio 7219 4414

[FeI/H] −0.19±0.05 −0.19±0.06 [FeII/H] −0.19±0.05 −0.17±0.06 [Na/Fe] −0.01±0.06 −0.05±0.05 [Mg/Fe] −0.07±0.05 0.11±0.06 [Al/Fe] ...... [Si/Fe] 0.23±0.04 0.16±0.07 [Ca/Fe] −0.50±0.09 −0.25±0.04 [ScII/Fe] 0.21±0.08 0.17±0.09 [Ti/Fe] −0.18±0.05 −0.03±0.06 [TiII/Fe] 0.08±0.08 0.01±0.05 [V/Fe] −0.11±0.09 0.14±0.06 [Ni/Fe] 0.04±0.06 −0.03±0.07 [Zr/Fe] −0.37±0.07 −0.42±0.05 [BaI/Fe] −0.05±0.00 −0.02±0.02 [LaII/Fe] −0.22±0.02 −0.14±0.04 [EuII/Fe] 0.08±0.01 0.07±0.03

Table B.12. Measured Abundances for Berkeley 85.

Ratio 4147 5002 5129

[FeI/H] −0.09±0.05 −0.12±0.05 −0.08±0.06 [FeII/H] −0.09±0.05 −0.09±0.05 −0.06±0.07 [Na/Fe] 0.25±0.06 0.21±0.05 0.05±0.07 [Mg/Fe] −0.19±0.03 0.03±0.00 −0.08±0.00 [Al/Fe] −0.04±0.01 0.04±0.07 −0.07±0.10 [Si/Fe] 0.15±0.05 0.11±0.02 0.11±0.07 [Ca/Fe] −0.16±0.07 −0.19±0.10 −0.23±0.10 [ScII/Fe] −0.02±0.05 −0.11±0.04 −0.04±0.07 [Ti/Fe] −0.02±0.07 0.01±0.03 −0.12±0.05 [TiII/Fe] 0.02±0.03 0.08±0.06 0.01±0.06 [V/Fe] −0.16±0.05 −0.22±0.06 0.01±0.11 [Ni/Fe] 0.00±0.06 0.00±0.03 −0.04±0.08 [Zr/Fe] −0.20±0.06 −0.39±0.05 −0.30±0.09 [BaII/Fe] 0.14±0.00 0.19±0.00 −0.02±0.00

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125 VITA Personal Background Julia Evans O’Connell Nashville, TN Daughter of Dudley Reed Evans and Mildred Lucille Gussow Children, Sally Marie and Daniel Joseph O’Connell

Education Diploma, McMinn County High School, Athens, TN, 1976 Bachelor of Science, Physics, Tennessee State University, Nashville, TN, 2011 Doctor of Philosophy, Physics, Texas Christian University, Fort Worth, TX, 2017

Experience REU research, Indiana University, Bloomington, IN, 2009 REU research, Cerro Tololo Inter-American Observatory, La Silla, Chile 2011 Teaching assistantship, Texas Christian University, Fort Worth, TX, 2012-2015, 2017 Research assistantship, Texas Christian University, Fort Worth, TX, 2015-2016

Professional Memberships American Astronomical Society ABSTRACT

SPECTROSCOPIC ANALYSES OF NEUTRON CAPTURE ELEMENTS IN OPEN CLUSTERS

by Julia E. O’Connell, Ph.D., 2017 Department of Physics and Astronomy Texas Christian University

Dissertation Adviser: Peter M. Frinchaboy III, Associate Professor of Physics and Astronomy

The evolution of elements as a function or age throughout the Milky Way disk provides strong constraints for galaxy evolution models, and on star formation epochs. In an effort to provide such constraints, we conducted an investigation into r- and s-process elemental abundances for a large sample of open clusters as part of an optical follow-up to the SDSS-

III/APOGEE-1 near infrared survey. To obtain data for neutron capture abundance analysis, we conducted a long-term observing campaign spanning three years (2013-2016) using the McDonald Observatory Otto Struve 2.1-meter telescope and Sandiford Cass

Echelle Spectrograph (SES, R(λ/∆λ) ∼60,000). The SES provides a wavelength range of

∼1400 A,˚ making it uniquely suited to investigate a number of other important chemical abundances as well as the neutron capture elements. For this study, we derive abundances for 18 elements covering four nucleosynthetic families– light, iron-peak, neutron capture and α-elements– for ∼30 open clusters within 6 kpc of the Sun with ages ranging from

∼80 Myr to ∼10 Gyr. Both equivalent width (EW) measurements and spectral synthesis methods were

employed to derive abundances for all elements. Initial estimates for model stellar

atmospheres– effective temperature and surface gravity– were provided by the APOGEE

data set, and then re-derived for our optical spectra by removing abundance trends as

a function of excitation potential and reduced width (EW/λ). With the exception of

Ba II and Zr I, abundance analyses for all neutron capture elements were performed by generating synthetic spectra from the new stellar parameters. In order to remove molec- ular contamination, or blending from nearby atomic features, the synthetic spectra were modeled by a best-fit Gaussian to the observed data. Nd II shows a slight enhancement in all cluster stars, while other neutron capture elements follow solar abundance trends. Ba

II shows a large cluster-to-cluster abundance spread, consistent with other open cluster abundance studies. From log(Age) ∼8.5, this large spread as a function of age appears to replicate the findings from an earlier, much debated study by D’Orazi et al.(2009) which found a linear trend of decreasing barium abundance with increasing age.