A Dissertation Entitled a Study of X–Ray Binary Populations in Star

A Dissertation entitled

A Study of X–ray Binary Populations in Star-Forming Galaxies

by Paula N. Johns Mulia

Submitted to the Graduate Faculty as partial fulﬁllment of the requirements for the Doctor of Philosophy Degree in Physics

Dr. Rupali Chandar, Committee Chair

Dr. Jon Bjorkman, Committee Member

Dr. S. Thomas Megeath, Committee Member

Dr. Scott Lee, Committee Member

Dr. Andrea Prestwich, Committee Member

Dr. Cyndee Gruden, Dean College of Graduate Studies

This document is copyrighted material. Under copyright law, no parts of this document may be reproduced without the expressed permission of the author. An Abstract of A Study of X–ray Binary Populations in Star-Forming Galaxies by Paula N. Johns Mulia

Submitted to the Graduate Faculty as partial fulﬁllment of the requirements for the Doctor of Philosophy Degree in Physics The University of Toledo December 2018

X–ray binaries (XRBs) are made up of a pair of closely orbiting stars, where one is a compact object (neutron star or black hole) that is accreting material from the other “donor” star. This material accumulates in an accretion disk, which can be heated to millions of degrees and emit large amounts of X–ray radiation. Depending on the mass of the donor star, XRBs are classiﬁed as either high–, intermediate–, or low–mass X–ray binaries (HMXBs, IMXBs, and LMXBs, respectively). Early- type galaxies only contain LMXBs, since no star formation has occurred for at least a billion years in these galaxies. Because of the smooth and easily modeled light distribution of early-type galaxies, their LMXB populations have been well-studied.

Far less is known about XRB’s in late-type, star-forming galaxies, which continue to form stars to this day, and therefore contain a mix of all three types. It is not possible to separate LMXBs, IMXBs, and HMXBs from X–ray properties alone, so researchers have resorted to making a number of assumptions and statistical corrections when studying XRBs in star-forming galaxies.

In this thesis, I develop two new methods for classifying XRBs from their optical properties, one based on the age of the parent star cluster and the other directly from the estimated mass of the donor star. We apply these to four star-forming galaxies: a high-intensity, on-going merger between two galaxies (the Antennae), a dwarf starburst (NGC 4449), and two typical spirals (M101 and M83). This is the iii first time XRBs have been classified on a source-by-source basis, and the preliminary results are intriguing. We find that IMXB properties appear to be more similar to

LMXBs than to HMXBs, and that HMXBs may have a diﬀerent formation mechanism than XRBs with lower mass donors. We discuss how this work can be expanded upon in the future to better characterize the properties of XRBs in star-forming galaxies.

iv For my husband– Alex, my conﬁdante and partner in all things. And for my grand- father, James Johns who always believed in me. Acknowledgments

I would ﬁrst like to thank my adviser, Rupali Chandar, for supporting me both inside and outside academia. Without her I would not have been able to grow and thrive in all aspects of my life while in graduate school.

To my family, without whom none of my success would be possible, especially to my grandparents, who showed through example how far hard work can take you, to my parents who inspired me to follow my dreams, and to my husband, who loved and supported me as a friend and as a partner all throughout graduate school.

I would lastly like to thank my friends and fellow graduate students. Graduate school was a marathon. I could not have crossed that ﬁnish line without you.

vi Contents

Abstract iii

Acknowledgments vi

Contents vii

List of Tables x

List of Figures xi

List of Abbreviations xiii

List of Symbols xiv

1 Introduction 1

1.1 X–ray Binaries ...... 1

1.2 Project Motivation ...... 6

1.3 Dissertation Outline ...... 10

2 Does High Density or Mass Help Star Clusters Produce X–ray Bi-

naries in Star–Forming Galaxies? 11

2.1 Introduction ...... 12

2.2 Data and Source Catalogs ...... 14

2.2.1 X–ray Source Catalogs ...... 14

2.2.2 Star Cluster Catalogs ...... 17

vii 2.2.3 Cluster Size and Density Estimates ...... 18

2.3 The Association between Clusters and XRBs ...... 21

2.4 Analysis ...... 25

2.4.1 Cluster Mass Distribution ...... 25

2.4.2 Cluster Size Distribution ...... 26

2.4.3 Cluster Density Distribution ...... 29

2.5 Discussion ...... 30

2.5.1 Dependence of XRB Production on Cluster Masses ...... 30

2.5.2 Dependence of XRB Production on Cluster Density ...... 33

2.5.3 Interpretation and Conclusions ...... 35

3 New Methods for Separating High and Low Mass X–ray Binaries in

Star–Forming Galaxies with an Application to M101 38

3.1 Introduction ...... 39

3.2 Catalog and X–ray Source Classiﬁcation ...... 43

3.2.1 Catalog of X–ray Binaries in M101 ...... 43

3.2.2 HST Observations and Alignment ...... 44

3.3 Methods for Classifying X–ray Binaries ...... 46

3.3.1 Classiﬁcation Based on Donor Star Masses ...... 46

3.3.2 Classiﬁcation Based on Parent Clusters Ages ...... 61

3.4 Spatial Distributions of XRBs ...... 63

3.5 X–ray Luminosity Functions ...... 64

3.6 Discussion and Conclusions ...... 71

4 Populations of X–ray Binaries in the Spiral Galaxy M83 72

4.1 Background ...... 72

4.2 Observations ...... 75

4.2.1 X–ray Observations with Chandra ...... 75

viii 4.2.2 HST Observations ...... 75

4.3 X–ray Binary Donor Catalogue ...... 76

4.3.1 Classiﬁcation Based on Donor Star Masses ...... 84

4.3.2 Classiﬁcation Based on Parent Cluster Ages ...... 87

4.3.3 Spatial Distribution of XRBs ...... 104

4.4 X–ray Luminosity Function ...... 104

4.5 Discussion ...... 110

5 Conclusions and Future Prospects 112

5.1 Conclusions ...... 112

5.2 Future work ...... 114

References 119

ix List of Tables

2.1 Parent Clusters to XRBs in the Antennae ...... 24

2.2 Parent Clusters to XRBs in NGC 4449 ...... 25

3.1 M101 X–ray Point Source Properties ...... 55

4.1 M83 X–ray Point Source Properties ...... 89

5.1 Galaxy Sample and Basic Properties ...... 116

x List of Figures

1-1 A diagram showing how X–ray emission is created in HMXBs and LMXBs. 2

1-2 Cumulative XLF of LMXBs in 14 E and S0 galaxies...... 5

1-3 Optical (HST) image of M101 with the locations of XRBs overplotted. . 8

1-4 Color-magnitude diagram of donor stars in M101...... 9

2-1 An inverted V-band image of the merging Antennae...... 15

2-2 An inverted V-band image of NGC 4449...... 16

2-3 Ishape FWHM versus C for clusters in the Antennae and NGC 4449. . . 20

2-4 Measurement of artiﬁcial cluster sizes in the Antennae and NGC 4449. . 22

2-5 The distribution of cluster masses in the Antennae and NGC 4449. . . . 27

2-6 Masses, sizes, and stellar densities for clusters in the Antennae and NGC

4449...... 28

2-7 The distribution of cluster densities in the Antennae and NGC 4449. . . 31

2-8 A comparison of the luminosity LX of XRBs coincident with a star cluster in the Antennae and NGC 4449 as a function of host cluster mass. . . . . 34

3-1 The cumulative X–ray luminosity functions for 29 nearby star-forming

galaxies, normalized by their star–formation rates...... 41

3-2 HST V–band mosaic of the 10 ﬁelds in M101...... 45

3-3 1.7”×1.7” thumbnails centered on X–ray point sources in M101. . . . . 46

3-4 Potential donor stars in M101 compared to theoretical mass tracks from

the Padova models at solar metallicity...... 52

xi 3-5 A comparison of measured host cluster B − V and V − I colors in M101

with predictions from the solar metallicity stellar evolutionary models of

Bruzual & Charlot (2003)...... 62

3-6 HST mosaic image showing the locations of classiﬁed X–ray sources in

M101 ...... 65

3-7 Cumulative X–ray luminosity functions for classiﬁed HMXBs, IMXBs, and

LMXBs in M101...... 66

3-8 Statistical ﬁts to total XLF of XRBs in M101...... 68

3-9 Statistical ﬁts to XLF of HMXBs, IMXBs, and LMXBs in M101. . . . . 70

4-1 A composite RGB image of M83 taken with the HST ...... 73

4-2 A composite RGB image of the nuclear region of M83 taken with the HST . 74

4-3 Locations of the seven M83 ﬁelds...... 77

4-4 1.7”×1.7” thumbnails centered on X–ray point sources in M83...... 77

4-5 Potential donor stars in M83 compared to theoretical mass tracks from

the Padova models at solar metallicity...... 86

4-6 The spatial distribution of X–ray sources in M83...... 105

4-7 Cumulative X–ray luminosity functions for classiﬁed HMXBs, IMXBs, and

LMXBs in M83...... 106

4-8 Statistical ﬁts to total XLF of XRBs in M83...... 107

4-9 Statistical ﬁts to XLF of HMXBs, IMXBs, and LMXBs in M83...... 109

xii List of Abbreviations

ACIS ...... Advanced CCD Imaging Spectrometer ACS ...... Advanced Camera for Surveys AGN ...... active galactic nucleus

C ...... concentration index CMD ...... color magnitude diagram

FWHM ...... full width at half maximum

HMXB ...... High–mass X–ray binary HLA ...... Hubble Legacy Archive HST ...... Hubble Space Telescope

IMF ...... initial mass function IMXB ...... Intermediate–mass X–ray binary

LMXB ...... Low–mass X–ray binary

PSF ...... point spread function

SNe ...... supernovae SNR ...... supernova remnant

ULX ...... Ultraluminous X-ray Source

WF ...... Wide Field CCDs WFC ...... Wide Field Camera WFC3 ...... Wide Field Camera 3 WFPC2 ...... Wide Field Planetary Camera 2

XLF ...... X–ray luminosity function XMM ...... X–ray Multi-Mirror Mission XRB ...... X–ray binary

xiii List of Symbols

α ...... exponent for the power law X–ray luminosity distribution β ...... exponent for the power law cluster mass distribution

βx ...... exponent for the power law simulated parent–cluster mass distribution ζ ...... exponent for the power law cluster density distribution

ζx ...... exponent for the power law simulated parent–cluster density distribution

Lx ...... X–ray luminosity L∗ ...... The X–ray luminosity value where a downturn in the Schechter function can be observed

N0 ...... A statistic which assesses whether or not there is a signiﬁcant downturn on the bright end of a distribution. A value >1 at a S/N level >5 suggests that there is

M ...... solar mass

MV ...... absolute V band magnitude Mv ...... apparent V band magnitude reff ...... cluster eﬀective or half–light radius ph ...... half–mass cluster density τ ...... age of a cluster population

xiv Chapter 1

Introduction

1.1 X–ray Binaries

X–ray binaries are two stars that are closely orbiting one another and which emit large amounts of X–ray radiation. In such a system, the more massive star will com- plete its hydrogen burning lifetime, explode as a supernova (without gravitationally unbinding its companion), and end its life as a neutron star or black hole. In low mass X–ray binaries (LMXBs), X–ray radiation is produced by the superheating of the accretion disk (fed by Roche-lobe overﬂow) around the compact object. In high mass X–ray binaries (HMXBs), X–rays are created when material from the donor star falls onto the compact object through strong stellar winds and/or by beginning atmospheric Roche-lobe overﬂow (Figure 1-1; Tauris & van den Heuvel 2006). In the case where the compact object is a blackhole, X–rays are produced by jets and the superheating of the accretion disk.

There are two main classes of X–ray binaries (XRBs), which depend solely on the mass of the donor star: high and low-mass X–ray binaries (HMXBs, LMXBs respect- fully). High-mass X–ray binaries have massive donor stars, often a supergiant or Be star. While the literature doesn’t provide a strict demarcation in donor star mass

1 Figure 1-1: A diagram showing how X–ray emission is produced by the superheating of the accretion disk in LMXBs (bottom) and when material from the donor star falls onto the compact object through strong stellar winds and/or by beginning atmospheric Roche-lobe overﬂow in HMXBs (top). Credit: Tauris & van den Heuvel 2006

2 between LMXBs and HMXBs, LMXBs are often considered to have donor masses

/1M (Fabbiano 2006). Early-type galaxies such as elliptical or lenticular galaxies have used up or lost most of their interstellar material (i.e., gas and dust), and therefore mostly contain older, low-mass stars. These galaxies only contain LMXBs since all donor stars have low masses. LMXBs in early-type galaxies reside either in ancient globular clusters or in the ﬁeld (Kim & Fabbiano 2004).

There has been a signiﬁcant amount of research on the properties of LMXBs in early-type galaxies. LMXBs can be found both in the ﬁeld and in globular clusters.

Between ∼20-70% of LMXBs are found in globular clusters, where this fraction de- pends on the globular speciﬁc frequency of the host galaxy (e.g., Angelini et al. 2001;

Jordan et al. 2004; Kundu et al. 2002, 2007; Sarazin et al. 2003; Humphrey & Buote

2008). Globular speciﬁc frequency is deﬁned as the number of clusters per unit galaxy luminosity which is also a function of galaxy morphological type. In general, ∼4% of the globular clusters in a given galaxy host LMXBs regardless of morphological type

(Kundu et al. 2002, Fabbiano 2006). It is interesting to note that while the fraction of globular clusters that host LMXBs is fairly constant across all studied galaxies, within each galaxy, the fraction of LMXBs that reside in globular clusters varies with galactic morphological type. This suggests that the formation eﬃciency of LMXBs is strongly correlated with globular cluster properties.

The X–ray luminosity function (XLF) is an important diagnostic tool that is used to understand LMXB populations in early-type galaxies. This distribution can be de-

dN −α 38 −1 scribed by a double power law, dL ∝ L , with a typical break at LX ∼5×10 ergs s (Figure 1-2). At low luminosities, ranging from a few 1037 to 5×1038ergs s−1, the XLF

is best ﬁt with a slope of 1.8±0.2, while at luminosities above the break the XLF is

steeper with a slope of 2.8±0.6 (Kim & Fabbiano 2004). Of particular interest is the

break detected in NGC 4697 at ∼2–5 1038erg s−1 (Sarazin, Irwin & Bregman 2000).

3 Because this break occurs near the Eddington limit for an accreting neutron star, this break could possibly mark the transition from neutron star to black hole. The existence of this break is controversial, as various other sources claim that it either doesn’t exist or show that the break disappears after preforming a completeness correction. (See Humphrey & Buote 2004; Kraft et al. 2001; Sivakoﬀ, Sarazin & Irwin

2003).

Previous work has suggested that the XLF for HMXBs have a near universal shape, which is a power law with a slope ∼1.6 (Grimm et al. 2003). When normalized by the star formation rate (SFR) of the host galaxy, the XLFs in star–forming galaxies appear to all have similar but not identical amplitudes. It has been suggested that the dispersion in the amplitude of the XLF distribution of HMXBs has a real physical origin, rather than resulting from errors and uncertainties (Mineo et al. 2012).

In elliptical galaxies, the X–ray luminosity function (XLF) scales with total galaxy mass (or luminosity). Additionally, Kim & Fabbianno (2004) ﬁnd a correlation between K-band luminosity (which is proportional to stellar mass) and integrated LMXB luminosity. Dynamical interactions of stars and compact objects in dense cluster cores are thought to be the primary formation mechanism for LMXBs in globular clusters

(Peacock et al. 2009; Jordan et al. 2007). Diﬀering XLFs of globular cluster and ﬁeld

37 −1 LMXBs at LX <5×10 erg s suggests they are two distinct populations (Maccarone et al. 2005; Irwin 2005; Juett 2005; Kim et al. 2009). Two plausible mechanisms have been suggested for the formation of ﬁeld LMXBs. One possibility is that they form in situ from evolved tight stellar binaries. However, because for a given stellar mass

LMXBs are predominantly found in globular clusters, the second possible formation mechanism is that LMXBs are formed exclusively in globular clusters and that the lighter ones are dynamically expelled from the cluster into the ﬁeld.

In elliptical galaxies, LMXBs are more likely to be found in red, younger and/or metal rich, clusters (V–I >1.1) (e.g., Kundu, Maccarone & Zepf 2002; Jordan et

4 Figure 1-2: Cumulative XLF of LMXBs in 14 E and S0 galaxies (Kim & Fabbiano 2004). The best ﬁt single power- law (dashed), and broken power-law (solid line) are shown. Note the marked location of the break at 38 −1 LX ∼5×10 ergs s .

5 al. 2004; Trudolyubov & Priedhorsky 2004; Kim et al. 2005). Although the cause of this correlation is not well understood, various explanations have been suggested

(e.g., Ivanova 2005; metallicity-dependent convective zone; Maccarone et al. 2004, irradiation-induced stellar winds).

X–ray intensity is typically measured in three wavelength ranges: the “soft” (S) band (0.3–1 keV), “medium” (M) band (1–2 keV), and “hard” (H) band (2–8 keV).

Two diﬀerent X–ray colors are typically deﬁned: H1 = (M - S)/T and H2 = (H -

M)/T, as well as a total X–ray luminosity (T=S+M+H). Kim et al. (2005) found no differences between the X–ray colors of LMXBs in globular clusters versus those found in the field, supporting the suggestion that all LMXBs form in globular clusters, but that some are expelled into the field.

We know a lot about the properties of XRBs in elliptical galaxies. In spiral, starburst, and merging galaxies however, recent star formation results in a complicated mix of LMXBs and HMXBs. As a result, far less is known about the XRBs in star–forming galaxies.

1.2 Project Motivation

In the past, two general methods have been used to separate HMXBs from LMXBs in star-forming galaxies. The ﬁrst attempts to use X–ray color to broadly classify X– ray sources into the diﬀerent groups based on where they lie on the X–ray color–color diagram (See Prestwich et al. 2003, 2009). Because HMXBs accrete material from their donor stars at a rate that is ∼100–10,000 times faster than LMXBs (Tauris & van den Heuvel 2006), HMXBs are typically associated with hard X–ray emission.

Soft X–rays are thought to originate from an optically thick accretion disk around a compact source, as luminous soft sources are therefore associated with LMXBs

6 while dim (< 1.0×1037 ergs s−1) soft sources are thought to be thermal supernovae remnants.

The second method simply excises XRBs in the parts of a galaxy believed to be dominated by LMXBs. For example, in M101, the bulge region is typically excluded because the X–ray population is thought to be dominated by LMXBs (see Figure 1-3); however a preliminary analysis by our group suggests that many of the X–ray sources in the “bulge” region of M101 are more likely to be HMXBs than LMXBs.

In this dissertation, we present two new methods for separating HMXBs, IMXBs, and LMXBs on a source-by-source basis. The deﬁnition of these populations is driven by the mass of the donor star, which cannot be uniquely determined from X–ray color, luminosity, or by association with a speciﬁc portion of a galaxy. Potential donor stars can however, be detected directly (or have an upper limit placed on their mass) from deep optical images, since stars emit much of their radiation at these wavelengths.

Some also continue to reside within the star cluster in which they were born, their parent cluster, and the age of the parent cluster provides important clues to the mass of the donor star.

The Hubble Space Telescope (HST) provides deep optical images of many nearby star-forming galaxies. Figure 1-4 demonstrates the potential for using deep, optical images to identify and directly determine the masses of donor stars. After matching the locations of X–ray point sources to optical images, X–ray sources that originate from foreground stars, hot, diffuse gas, and background AGN can be excluded, leaving a catalogue of XRBs in the host galaxy. Potential donor star(s) are identified by matching optical sources within a 2σ positional uncertainty to each X–ray point source. Multi-band photometry of each donor star gives their luminosities and colors, which are compared with theoretical isochrones to estimate its mass. The HST data are sufficiently deep that we should be able to detect donor stars down to ∼3 M , i.e. very close to the LMXB/HMXB separation limit.

7 Figure 1-3: Optical (HST) image of M101 with the locations of XRBs overplotted (From Mineo et al. 2012). The crosses indicate the positions of detected X–ray point sources used for analysis. The red circle marks the estimated radius of the bulge. X–ray sources within this radius are marked by small red circles and were excluded from the analysis because it was assumed that they are dominated by LMXBs (see Gilfanov 2004, for details).

8 Figure 1-4: Color-magnitude diagram of donor stars in M101.

9 Recent works have identiﬁed a number of X–ray binaries in star clusters (Rangelov et al. 2012, Fall et al. 2005, Zezas et al. 2002; Kaaret et al. 2004). In these cases, no donor star can be identiﬁed because of the severe crowding. However, multi-band photometry of a cluster provides an estimate of its age. If the cluster is younger than

10 Myr, the XRB is very likely to have a high mass donor star (>8 M ), since the most massive stars in a cluster are the most dynamically active.

1.3 Dissertation Outline

The rest of this dissertation is organized as follows. In Chapter 2 we identify the parent clusters to XRBs in the merging Antennae galaxies and in the starburst

NGC 4449, and use their properties to assess whether higher cluster mass or density plays a role in the formation of XRBs. In Chapter 3, we present our new methodology for classifying XRBs based on the mass of the donor star, and apply this method to sources in M101. In Chapter 4, we extend the application of this method to XRBs in another nearby spiral galaxy, M83, in order to take a deeper look at XRB formation and evolution theories. Finally, in Chapter 5, we summarize our main conclusions and brieﬂy discuss some directions for future work.

10 Chapter 2

Does High Density or Mass Help

Star Clusters Produce High Mass

X–ray Binaries in Star–Forming

Galaxies?∗

It is now widely accepted that many X–ray binaries form in compact star clusters.

Therefore, it follows that there may be a link between the properties of star clusters

and the production of X–ray Binaries (XRBs). We study the masses and densities

of compact clusters with diﬀerent ages found in the merging Antennae galaxies and

in the dwarf starburst NGC 4449, to assess their impact on the production of XRBs.

The very youngest clusters, with ages log (τ/yr) < 7.0 likely host high-mass XRBs, while older clusters with ages log (τ/yr) = 7.0 − 8.6 host intermediate-mass XRBs.

We ﬁnd tentative evidence that intermediate-mass XRBs are more likely to form in more massive and dense clusters, while high-mass XRBs do not appear to form in

∗Paula Johns Mulia, Rupali Chandar, Blagoy Rangelov Accepted for publication to the Astrophysical Journal

11 clusters that favor higher masses or densities. If conﬁrmed, these ﬁndings support a

picture where high- and intermediate-mass XRBs form through diﬀerent mechanisms.

2.1 Introduction

X–ray observations of star forming galaxies such as the merging Antennae and the

dwarf starburst NGC 4449 taken with the Chandra X–ray Observatory (Chandra)

show many point sources, believed to be X–ray binaries (XRBs) (Zezas et al. 2002;

Zezas et al. 2006; Rangelov et al. 2011). There are three broad classes of XRBs,

which depend on the mass of the donor star: high-, intermediate-, and low-mass X–

ray binaries (HMXBs, IMXBs, and LMXBs, respectively). The literature does not

provide consistent deﬁnitions for the mass ranges of these donor stars, so in this work

we consider a HMXB to have a donor star with a mass ≥ 8 M , an IMXB to have a

donor mass of 3 − 8 M , and a LMXB to have a donor mass < 3 M . Although many XRBs are found close to star clusters, some continue to reside

within their parent cluster (e.g., Zezas et al. 2002; Kaaret et al. 2004; Rangelov et

al. 2012). In these cases, the age of the parent cluster provides the age of the binary.

Age also provides strong clues to the type of XRB, since the most massive remaining

cluster stars are the most dynamically active and hence have a high probability of

being the donor. Massive ≥ 8 M stars have hydrogen burning lifetimes of ≈ 10 Myr,

while 3 M stars have lifetimes of ≈ 300 − 400 Myr. This suggests that clusters with ages ≤ 10 Myr are most likely to host HMXBs, those with ages between ≈ 10–400 Myr to host IMXBs, and clusters older than 400 Myr host LMXBs.

In early-type galaxies, observations and simulations suggest that LMXBs in ancient globular clusters form through dynamical interactions. N-body simulations ﬁnd that the high rate of interactions in dense stellar cores can easily form new binaries or tighten existing ones. Observations also indicate that higher density globular clus-

12 ters are more likely to host a LMXB. For example, Sivakoﬀ et al. (2007) found that

globular clusters with smaller radii are more likely to form LMXBs. Because clus-

ter size and mass are unrelated (Jord´anet al. 2005), this implies that the globular

clusters that host LMXBs tend to be more dense. Additionally, LMXBs in elliptical

galaxies are more likely to be found in red globular clusters (V−I > 1.1), which tend to be more centrally concentrated (and hence more dense) than their blue counterparts (e.g., Kundu, Maccarone & Zepf 2002; Jordan et al. 2004; Trudolyubov &

Priedhorsky 2004; Kim et al. 2005, Peacock et al. 2016).

Far less is known about the formation of XRBs in star-forming galaxies, which contain HMXBs and IMXBs in addition to LMXBs. While dynamical interactions in the dense cores of young clusters may be responsible for the formation of HMXBs and

IMXBs, it is also possible that primordial pairs of high mass stars are the progenitors in young stellar systems (Ivanova et al. 2008; Garofali et al. 2012). In this case, after the more massive star ﬁnishes its life and ends as a compact object (neutron star or black hole) and if it remains as a tight binary, mass could be transferred to the compact object from the remaining massive (donor) star.

To date, there have been very few studies testing whether the mass and/or density of clusters younger than few × 100 Myr aﬀects the production rate of XRBs, because of the complicated morphologies in the star-forming galaxies where these young systems form. Clark et al. (2008) made an early test of the relationship between HMXB frequency and total cluster mass in the Antennae; however their results were inconclu- sive, because they used a crude age–dating method, forcing them to make a number of assumptions.

In this work, we investigate whether the masses and/or densities of massive young– star clusters in the merging Antennae and the dwarf starburst NGC 4449 aﬀects the rate at which XRBs are produced. The rest of this paper is organized as follows;

Section §2 summarizes the X–ray and optical observations of the Antennae galaxies

13 and NGC 4449 from Chandra and Hubble Space Telescope (HST ) used in this work.

Section §3 describes the astrometric matching between the X–ray and optical source catalogs. §4 presents the overall distribution of cluster masses and densities, and compares with those that have produced HMXBs. A discussion of results and the implication for the formation of XRBs in clusters is found in §5.

2.2 Data and Source Catalogs

2.2.1 X–ray Source Catalogs

+0.15 Antennae: We assume a distance of 21 ± 2.6 Mpc (distance modulus= 31.71−0.41 mag) to the Antennae (Schweizer et al. 2008; Riess et al. 2011). We use the X–ray point source catalog published by Zezas et al. (2006), based on seven exposures taken with the ACIS instrument on Chandra, which includes positions and total X–ray luminosities. Here, we adopt the average X–ray luminosity from the 6 observations. Zezas et al. (2006) eliminated foreground stars, detections of diﬀuse emission, and background AGN, leaving a total of 82 X–ray sources in their catalog. The locations of these sources are superimposed on an optical image of the Antennae in Figure 2-1.

NGC 4449: We assume a distance of 3.8 ± 0.27 Mpc (distance modulus of

27.9 ± 0.15) to NGC 4449. We use the X–ray point source catalog published by

Rangelov et al. (2011), based on 3 sets of Chandra observations taken with the ACIS instrument. This catalog includes the positions and X–ray luminosities of 23 X–ray point sources in NGC 4449. The locations of these sources are superimposed on an optical image of NGC 4449 in Figure 2-2.

14 Antennae

5’’

Figure 2-1: An inverted V-band image of the merging Antennae. Red dots represent the location of star clusters while blue crosses represent candidate X–ray binaries. At the 21 Mpc distance of the Antennae, 100 = 101.8 pc.

15 NGC 4449

10”

Figure 2-2: An inverted V-band image of NGC 4449. Red dots represent the location of star clusters while blue crosses represent candidate X–ray binaries. At the 3.8 Mpc distance of NGC 4449, 100 = 18.4 pc.

16 2.2.2 Star Cluster Catalogs

Antennae: We use the cluster catalog from Whitmore et al. (2010), which was ﬁne–tuned by Rangelov et al. (2012). Clusters were selected to be point-like sources brighter than MV = −9, since few individual stars reach these luminosities. The Rangelov et al. (2012) catalog includes the locations, photometry, mass and age estimate for ≈ 700 clusters. Figure 2-1 shows an inverted V-band image of the

Antennae overlaid with the location of X–ray point sources (blue crosses) and clusters

(red dots).

The age and mass of each cluster, was determined in Whitmore et al. (2010) by comparing the observed UBVI & Hα magnitudes with those predicted by the population synthesis models of Bruzual & Charlot (2003). They assumed a solar metallicity, a Chabrier initial mass function, and a Galactic-type extinction curve. The age-dependent mass-to-light ratios from the models were used with the extinction- corrected V-band luminosity to estimate the mass of each cluster (Whitmore et al.

2010). The dominant source of uncertainty in cluster masses comes from uncertainties in the age estimates, because errors in mass and age are correlated. Whitmore et al.

(2010) estimated these errors to approximately a factor of 2, or 0.3 in log M. Uncer- tainty in the distance only aﬀects the cluster mass estimates at the ≈ 10% level, and photometric uncertainties have an even smaller impact, because these are all fairly bright sources.

NGC 4449: We use the cluster catalog published by Rangelov et al. (2011), who selected clusters to be broader than the PSF. Their catalog includes the coordinates, photometry, mass and age estimates for 129 clusters. They also estimated the sizes, reff , of each cluster. Figure 2-2 shows the location of each candidate cluster and superimposed on a V-band image of NGC 4449 (red circles). The mass and age estimates were made using a similar method as that used for clusters in the Antennae,

17 and the masses are uncertain at a similar factor of two level. Uncertainties in the distance aﬀect the mass estimates for clusters in NGC 4449 at the ≈ 10% level, and photometric errors have an even smaller impact.

2.2.3 Cluster Size and Density Estimates

The sizes of clusters in NGC 4449 were measured by Rangelov et al. (2011), but no such measurements have been published for clusters in the Antennae. Here, we follow a similar procedure as Rangelov et al. (2011) to estimate the half-light radii of clusters in the Antennae, so that the sizes (and densities) of clusters in the two systems can be compared directly.

We use the BAOLAB/Ishape software (Larsen 1999) to estimate the half light radius or reff for each cluster in the Antennae. Ishape fits analytic profiles convolved with the point spread function (PSF) to determine the best fit radius of each cluster.

We create a PSF from '20 relatively bright, isolated stars. We assume a King proﬁle

(King 1966) with a ratio of tidal to core radius of 30 to represent the typical cluster density light proﬁle, although the results are not sensitive to the exact value of rt/rc that is assumed. Based on extensive prior testing of ISHAPE, typically any source with a FWHM measurement of 0.2 pixels or less is considered to be indistinguishable from the PSF (e.g., Larsen 1999; Chandar et al. 2016). Because the Antennae is ∼5 times further away than NGC 4449, cluster sizes are more diﬃcult to measure. Many of the youngest (τ < 10 Myr) clusters are quite compact and in crowded regions, and cannot easily be distinguished from point sources at this distance.

We also measure the concentration index (C ) for each cluster, where C is the diﬀerence in magnitude measured in a 3 pixel and 0.5 pixel aperture, for each cluster.

C is a crude but fairly stable estimate of cluster size. (Figure 2-3) shows that FWHM and concentration index measurements are fairly well correlated for clusters in the

Antennae (ﬁlled symbols) and NGC 4449 (open symbols). The relation is steeper at

18 smaller radii and ﬂattens at larger ones, because the concentration index ‘saturates’,

since the largest aperture size is only 3 pixels.

The measured FWHM is converted from pixels to reff by using a pixel scale of 0.0005 pixel−1, an assumed distance of 21 Mpc to the Antennae, and a conversion factor of 1.48 as indicated in the Ishape manual (Larsen 1999). Since sources with sizes of 0.2 pixels and less are considered indistinguishable from a point source, we cannot measure reff values lower than ∼ 1.5 pc at the distance of the Antennae. At the distance of NGC 4449 this size limit of 0.2 pixels corresponds to ≈ 0.3 pc. We ﬁnd

a median size of reff ∼ 2.5 pc (reff ∼ 2 pc) for clusters in the Antennae (NGC 4449).

Not surprisingly, these are somewhat smaller than the median size reff ≈ 3 − 4 pc found for Galactic globular clusters, because young clusters tend to be more compact

than older ones.

As a check on the accuracy of the Ishape size measurements, we create artiﬁcial

clusters with known sizes, add them to each image, and then measure their FWHM

with Ishape. Speciﬁcally, we generate 8,000 artiﬁcial clusters for each galaxy with:

(1) FWHM= 1.0 pixels, (2) the median cluster size in each galaxy, FWHM∼ 0.3 pix

(∼ 2.5 pc) for the Antennae and ∼ 1.5 pixels (∼ 2 pc) for NGC 4449, and (3) a a common size of 4 pc, which translates to 0.5 pix (1.9 pix) for the Antennae

(NGC 4449). The artiﬁcial clusters were generated using the MKCMPPSF task in

BAOLAB, which convolves a PSF with a KING30 proﬁle. The MKSYNTH task added artiﬁcial clusters to each image with magnitude ranges matching the observed range in each galaxy.

Figure 2-4 shows the results of our artificial cluster experiments. These figures plot the input versus measured reff for each artificial cluster as a function of absolute V magnitude. The horizontal lines are the input sizes of the artificial clusters. Overall, for the full range of cluster magnitudes studied in this work, the experiments do a very good job of measuring cluster size. We find that at least 95% (98%) of the

19 Figure 2-3: This figure shows Ishape FWHM versus C for clusters in the Antennae (filled) and NGC 4449 (hollow). Black filled circles represent clusters with associated HMXBs. In NGC 4449, two additional sources are coincident within 80pc, and are shown with crossed black circles. In the Antennae, clusters with FWHM<2 pix are fitted with linear relationship (dark line). Clusters in the An- tennae which fall to the left the dashed vertical line at 0.2 pixels are omitted from further size-related analysis. Clusters in NGC 4449 (hollow) can also be fitted with a simple linear relationship (see text).

20 artiﬁcial clusters have recovered sizes measured within 20% (50%) of their input sizes

Artiﬁcial clusters in crowded regions are somewhat more likely to have measured sizes that deviate from their actual size by 50%, but overall, the size estimates appear to be quite robust, likely because the clusters in both galaxies are fairly bright. Based on these experiments, we estimate a typical uncertainty of ≈ 20% in the size estimates of our clusters.

We also determine the stellar density of each cluster. Ignoring any internal

3 mass/light segregation, the half-mass density is deﬁned by ρh ≡ 3M/8πrh where we utilize the standard 3D-2D conversion rh =(4/3)reff . Clusters in the Antennae that

3 are younger than log(τ/yr)<7 have a median stellar density ∼ 80 M /pc , and those

3 with ages log(τ/yr)=7–8.6 have a median density of ∼ 600 M /pc . In NGC 4449,

3 clusters younger than log(τ/yr)≤7 have a median density ∼ 80 M /pc , and those

3 with ages log(τ/yr)=7–8.6 have a median density of ∼ 20 M /pc . Clusters in the An- tennae at all ages tend to have higher densities than their counterparts in NGC 4449, because they have significantly higher masses, rather than significant differences in their sizes.

2.3 The Association between Clusters and XRBs

Point-like sources detected in HST images taken with the ACS/WFC camera have astrometric accuracies better than ≈0.0005, while X–ray point sources detected in Chandra images taken with the ACIS-S instrument have astrometric uncertainties which are approximately an order of magnitude larger. Rangelov et al. (2012) estimated the positional uncertainty for X–ray sources in the Antennae by comparing the locations of X–ray point sources (e.g., background AGN & foreground stars) with that of their optical counterparts, located in BVI color images available from the

21 Figure 2-4: Measurement of artificial cluster sizes in the Antennae (top) and NGC 4449 (bottom). The input sizes are listed, and include the median cluster size found for each galaxy. Over the observed range of magnitudes, we find that ≈ 95% of artificial clusters have size measurements within 20% of their input size.

22 Hubble Legacy Archive (HLA)1. Using this procedure, the 1σ positional uncertainty for the X–ray observations is found to be ∼ 0.004 for the Antennae (Rangelov et al.

2012), and ∼0.005 for NGC 4449 (Rangelov et al. 2011).

Catalogs of both X–ray point sources and of star clusters have been previously published for the Antennae and NGC 4449. In the Antennae, at a distance of ∼21

Mpc, 9 XRBs are spatially coincident within 1σ (which translates to ∼40 pc) with a star cluster, while 23 XRBs are spatially coincident within 2σ positional uncertainty

(∼80 pc) with a star cluster. For NGC 4449, which is much closer at a distance of only 3.8 Mpc, 3 XRBs are spatially coincident within 1σ (∼9 pc) positional uncertainty of its candidate parent–cluster, while 4 XRBs are spatially coincident within

2σ (∼18 pc). If we consider the physical scale of a 2σ coincidence in the Antennae for NGC 4449, we ﬁnd two additional XRBs in NGC 4449 that are located within

80 pc of a star cluster.

Table 2.1 compiles properties of the 23 clusters (including age and mass) that are coincident, within 2σ of an X–ray point source in the Antennae, and Table 2.2 has similar properties for the 6 XRBs in NGC 4449 that are located within 2σ a star cluster. We consider these to be the parent cluster to each XRB.

1http://hla.stsci.edu

23 Table 2.1: Parent Clusters to XRBs in the Antennae

X–ray Source Host Cluster a X–ID log(LX ) Mass Age reff −1 (ergs s ) (log(M/M )) (log(τ/yr)) (pc) 11 38.3 1.3×105 8.36 1.8 22 37.5 4.5×104 6.56 < 1.5 24 37.5b 3.3×104 6.94 < 1.5 27 39.8 1.0×106 8.06 2.0 36 38.0 6.6×104 6.58 2.0 41 37.8 3.1×105 6.00 17.0 45 37.4 1.1×105 7.72 7.8 47 38.4 3.5×104 6.68 0.6 49 38.2 7.8×104 7.96 2.5 50 38.7 7.9×104 6.70 < 1.5 52 37.4b 6.6×105 7.51 2.7 53 38.2 8.0×104 6.22 < 1.5 58 38.4 1.1×105 6.84 < 1.5 83 38.6 1.1×105 8.06 2.9 84 38.8 2.1×105 6.50 17.6 87 37.8 4.1×104 6.74 < 1.5 89 37.3 4.3×104 6.38 18.3 91 37.4 1.6×106 8.31 3.1 98 37.6 4.2×105 7.60 1.6 99 39.5 2.0×105 8.21 1.8 101 38.0 5.0×105 6.00 0.96 121 37.5 1.6×105 6.54 7.1 a X–ID from Zezas et al. (2006) b LX could not be estimated. Instead, these values are taken from XRBs with similar counts. Note: LX is the average luminosity in the six observations taken by Zezas et al. (2006). Coincident XRBs are within 2σ (∼ 80pc) of their parent cluster.

24 Table 2.2: Parent Clusters to XRBs in NGC 4449

X–ray Source Host Cluster

a X–ID log(LX ) Mass Age reff

−1 (ergs s ) (log(M/M )) (log(τ/yr)) (pc) 12 38.5 2.2×104 6.96 1.0

15 37.9 1.7×104 8.31 4.3

24b 37.8 1.5×103 6.84 3.5

21 37.1 8.4×104 8.21 1.6

43b 36.5 5.2×103 6.40 7.9

46 36.9 7.3×103 6.00 0.7

a X–ID from Rangelov et al. (2012) b Coincident within 80pc Note: LX were derived by Rangelov et al. (2011). Coincident XRBs are within 2σ (∼18pc) of their parent cluster. Two additional sources are coincident within 80pc, and are discussed in §5.3. 2.4 Analysis

In this section, we characterize the mass and density distributions of the overall cluster populations in the Antennae and NGC 4449, and use statistical tests to assess whether or not clusters that host XRBs are drawn from these distributions or are weighted to higher masses and densities.

2.4.1 Cluster Mass Distribution

The mass function of young cluster populations in star-forming galaxies can be approximated by a power law distribution, dN/dM ∝ M β with β ∼ −2, including in the

25 Antennae (e.g., Zhang & Fall 1999; Whitmore et al. 2010) and NGC 4449 (Rangelov

et al. 2011). Speciﬁcally, the mass functions for clusters in the Antennae are best

ﬁt by β = −2.29±0.09 for log(τ/yr)≤7, and β = −1.93±0.07 for log(τ/yr)=7–8.6.

In NGC 4449, the cluster mass functions have best–ﬁt values of β = −1.83±0.20 for log(τ/yr)≤7, and β = −1.84±0.21 for log(τ/yr)=7–8.6 (Fig. 2-5).

The top panels of Figure 2-6 show the distribution of cluster masses as a function of age for the Antennae and NGC 4449. Black stars represent all selected clusters in the Antennae (left) and in NGC 4449 (right), and red diamonds represent potential parent–clusters to XRBs, as described in Section 3. Unﬁlled red diamonds represent the two additional XRBs that lie within ∼80pc of a cluster in NGC 4449. One important result is that the masses of clusters in NGC 4449 which host XRBs are all lower than their counterparts in the Antennae, well outside any uncertainties in the mass estimates.

This diﬀerence in cluster masses between the two star-forming systems is expected of populations with diﬀerent sizes that are drawn from a similar underlying power-law distribution. Galaxies with higher rates of star formation are expected to form more clusters and also clusters with higher masses. The Antennae has a star formation rate

−1 −1 (SFR) of ≈ 20 M yr , while NGC 4449 has a SFR of ≈ 0.35 M yr (Chandar et al. 2017), and the most massive diﬀer by a factor of ≈ 50 between the two galaxies.

We perform statistical tests to quantify if XRBs tend to form in more massive clusters in Section 5.1.

2.4.2 Cluster Size Distribution

We know less about the distribution of cluster sizes than about their masses.

There is at least some evidence that clusters expand rapidly, over the ﬁrst ∼ 10 Myr

(Bastian et al. 2012; Chandar et al. 2016). The center-left panel of Figure 2-6 shows the reff measured for clusters in the Antennae as a function of age (black stars).

26 Figure 2-5: The distribution of cluster masses in the Antennae (top) and NGC 4449 (bottom). The best ﬁt power law, dN/dM ∝ M β, are shown as the dashed lines (see Section 2.4.1).

27 Figure 2-6: Masses, sizes, and stellar densities for clusters in the Antennae (left) and NGC 4449 (right) are shown as the ﬁlled circles. Filled red diamonds show the properties of clusters that host XRBs, and the unﬁlled red diamonds represent the two additional clusters in NGC 4449 that are within 80 pc of a XRB (see text). A number of very young clusters in the Antennae are too compact to be distinguished from a point source, and only upper limits to their reff are shown (and hence lower limits to their densities).

28 Many of the young clusters do not appear to be broader than the PSF, and therefore

have upper limits to their size estimates of reff < 1.5 pc, including a number of XRB-hosts (red diamonds). All of the older XRB host clusters are resolved.

Clusters in NGC 4449 have reff ranging from ∼ 0.4–10 pc. The youngest of these clusters (log(τ/yr)≤7) are relatively compact and have a median size of ∼ 1.3 pc,

while intermediate age clusters (log(τ/yr)=7–8.6) have a median size of reff of ∼ 3.7 pc. The center-right panel of Figure 2-6 shows the distribution of cluster sizes in

NGC 4449 as a function of cluster age. There is an increasing trend of reff with age for the clusters in NGC 4449, similar to what is observed for M51 (Chandar et al.

2016), suggesting that clusters expand over the ﬁrst few 100 Myrs. Red diamonds show the properties of candidate parent–clusters to XRBs. Unﬁlled red diamonds represent the two additional clusters in NGC 4449 that lie within ∼80pc of an XRB.

2.4.3 Cluster Density Distribution

The ﬁlled circles in the bottom panels in Figure 2-6 show the estimated stellar

−3 densities (in M pc as determined in Section 2.3), of all clusters younger than ≈ 400 Myr in the Antennae (left) and NGC 4449 (right). The red diamonds show the density of parent clusters to XRBs, including the clusters which have lower limits to their estimated densities because we could only place an upper limit on their reff measurement.

The half-mass densities calculated here have a range similar to that of the half- mass densities found for galactic globular clusters and for YMCs in other galaxies

(e.g., Lee et al. 2005; McLaughlin & Fall (2008). Globular clusters in the Milky Way

4 3 Galaxy have half-mass densities ranging from 0.03-5.65×10 M /pc with a median

3 density of 246 M /pc (McLaughlin & Fall 2008). Figure 2-7 shows the cluster-density distributions in diﬀerent intervals of age.

These can be approximated by a power law, dN/dρ ∝ ρζ . The best ﬁt power-law

29 index ζ and formal uncertainties are given in each panel. Overall, the shape of the cluster density distributions are quite similar in both galaxies in the diﬀerent intervals of age, with ζ ≈ −1.0 ± 0.15.

2.5 Discussion

2.5.1 Dependence of XRB Production on Cluster Masses

To test the dependence of XRB production on cluster mass, we apply a one–tail t–test to our sample. This test compares the means of two groups to assess whether they are statistically diﬀerent from one another, and can be used when there are small numbers. This test is particularly useful here because we do not have enough XRB parent–clusters to determine a meaningful mass distribution. Instead, we compare the observed mean parent–cluster mass to the mean “simulated” parent–cluster mass

(described below), and use the t–test to determine if the simulated distribution (βx) can be ruled out. The p–value represents the probability that the observed parent– cluster masses are culled from the simulated parent–cluster mass distribution (βx).

For the purposes of this statistical test, distributions (βx) which return p < 0.03 are rejected.

We randomly generate n parent–cluster masses, where n is the number of observed parent–clusters, which follow the mass distribution we want to test, from the observed cluster mass range. We explore simulated parent–cluster mass distributions that are drawn from power laws with βx = -3 to -0.5, and for each βx we run the simulation 1,000 times and take the mean of each distribution. We take the mean simulated parent–cluster mass to be the median value of the 1,000 means. This method is used to rule out parent–cluster mass distributions that don’t match the observed parent– cluster sample to a statistically signiﬁcant degree. If the simulated mean is suﬃciently small or large in comparison to the observed mean parent cluster mass, the t-test will

30 Figure 2-7: The distribution of cluster densities in the Antennae (top) and NGC 4449 (bottom). The best ﬁt power law, dN/dρ ∝ ρζ , are shown as the dashed lines (see Sec- tion 2.4.3).

31 return a low probability (p < 0.03) that the XRB-parent population is drawn from

that particular simulated distribution.

In order to assess the role played by cluster mass in producing XRBs, we compare

simulated parent-cluster mass distributions (βx) to the mass distribution of the galaxy (β) for three diﬀerent age bins. We consider the cases where the parent–cluster mass

distribution of XRBs (βx) follow: (1) the overall cluster mass distribution of the galaxy (βx = β), (2) a power law with steeper (βx < β) or shallower slopes (βx > β), and (3) a uniform distribution (βx = 0). For the Antennae galaxies, our results at all ages rule out parent–clusters that have a mass distribution that is uniform (βx = 0) or steeper (βx < β) than that of the total observed cluster population. We ﬁnd that somewhat older (log(τ/yr) between 7.0 and 8.6) parent clusters in the Antennae are consistent with being drawn from a mass function weighted towards higher masses, while young parent clusters

(log(τ/yr)<7) are consistent with being drawn from the overall mass distribution of the cluster population. Although tentative because of the poor statistics, this is an intriguing result, which we further discuss below. Unfortunately, our parent–cluster population is too small for signiﬁcant statistical testing in NGC 4449.

One potential complication is that we cannot diﬀerentiate whether a given cluster has formed only one or multiple HMXBs. Is it possible that the most massive clusters host multiple HMXBs? Although we cannot deﬁnitively answer this question, a comparison of the X–ray luminosity LX of coincident HMXBs as a function of host cluster mass (Figure 2-8) shows that the X–ray emission is not obviously brighter in the most massive clusters than in lower mass ones, as we might naively expect if multiple XRBs form in more massive stellar systems. In addition, the two most

39 −1 luminous x-ray sources with LX >10 ergs s (ULXs), exhibit long term variability which suggests that they are dominated by a single luminous accreting source rather

32 than multiple sources. Several others with lower luminosities also exhibit long term

variability (Zezas et al. 2006).

We conclude that XRBs with ages of log(τ/yr)=7–8.6 may be more likely to form

in more massive clusters, but that XRBs in the youngest, log(τ/yr) < 7 clusters

appear to follow the underlying cluster mass distribution. These results suggests that

HMXBs, which form in the youngest clusters (τ ≤ 10 Myr), may have a diﬀerent formation mechanism than IMXBs, which are the most likely XRB contribution in clusters with ages between ≈ 10 and 400 Myr.

2.5.2 Dependence of XRB Production on Cluster Density

Previous works have found that higher density globular clusters in early-type galaxies are more likely to host LMXBs, because more frequent dynamical interactions can potentially lead to enhanced formation of tight binaries (e.g., Peacock et al. 2009;

Jordan et al. 2007). To test the role that cluster density plays on the production of

XRBs in star-forming galaxies, we apply a one–tail t–test to our sample, following a procedure similar to that laid out in the previous section. We consider the cases where the parent–cluster density distribution of XRBs (characterized by the power law index

ζx) follows: (1) the overall cluster density distribution of the galaxy (ζx = ζ), (2) a

power law with steeper (ζx < ζ) or shallower slopes (ζx > ζ), and (3) a uniform

distribution (ζx = 0). We ﬁrst present the results for older log (τ/yr) = 7.0 − 8.6 clusters, which likely

host an IMXB. Steep power laws with ζx < −1.0 have a low probability (/ 3%) of representing the density distributions of the XRB parent–clusters in the Antennae

at these ages. Instead, these parent clusters are consistent with being drawn from

higher densities than the overall distribution, although we cannot completely rule out

that they may be drawn from the underlying density distribution. This result, while

tentative, supports a scenario where dynamical interactions may play an important

33 Figure 2-8: A comparison of the luminosity LX of XRBs coincident with a star cluster in the Antennae (ﬁlled) and NGC 4449 (hollow) as a function of host cluster mass. The blue triangles and red squares represent clusters with log(τ) ≤ 7 and log(τ/yr)=7–8.6 respectively. In NGC 4449, two additional sources are coincident within 80pc, and are shown with crossed black circles. The two crossed–out sources correspond to X24 and X52, whose LX could not be estimated and instead were taken to be the same as XRBs with similar counts.

34 role in creating IMXBs in star-forming galaxies, just as they do for LMXBs in early-

type galaxies. The parent–cluster population in NGC 4449 for this age range is too

small for statistical testing.

Next, we consider the stellar densities of very young log (τ/yr) < 7.0 clusters.

In NGC 4449, the t-test results suggest that it is unlikely that the density distri-

bution of the log(τ/yr)<7 parent–clusters is steeper than a power law with a slope of ζx = −1.5 or that it follows a uniform distribution. We ﬁnd that the densities of log(τ/yr)<7 parent–clusters in NGC 4449 are consistent with being drawn from the density distribution of the general population. We are unable to compare the properties of XRB-host clusters in the Antennae with the overall population in this age interval, because many of them only have a lower limit to the density.

While tentative, our results for the masses and densities of XRB host clusters paint a consistent picture. For the youngest clusters, with ages log (τ/yr) < 7.0, where the

X–ray binary is likely to have a high mass donor, our results suggest that there is no tendency for HMXBs to form in clusters with higher mass (Antennae) or density

(NGC 4449). A diﬀerent result is found for clusters with ages log (τ/yr) = 7.0 − 8.6, which likely host an intermediate mass donor. Here, the IMXB host clusters tend to have higher mass and density than the underlying cluster distribution.

2.5.3 Interpretation and Conclusions

The parent clusters which produce HMXBs do not appear to require higher masses or densities than the overall cluster population at similar ages. This suggests that high stellar mass and density, and hence dynamical interactions, are not a primary driver of HMXB production. Our result is consistent with those from simulations which

ﬁnd that at these very young ages and typical densities, there isn’t suﬃcient time for

HMXBs to form via dynamical interactions (e.g., Garofali et al. 2012). Therefore, it

35 is likely that HMXBs form from a primordial pair of high mass stars (e.g., Ivanova et

al. 2008).

The parent clusters which produce IMXBs do appear to favor higher masses and

densities than the overall cluster population at similar ages. This result suggests

that dynamical interactions play an important role in the production of IMXBs in

stellar clusters. Dynamical interactions have also been found to be important in the

production of LMXBs in the globular cluster systems of early-type galaxies, and so

the formation of IMXBs may be more similar to that of LMXBs than of HMXBs.

Our work shows that XRBs can form in star clusters that have a large range of

stellar masses. We ﬁnd XRBs in the Antennae and NGC 4449 that have formed with

overlapping X–ray luminosities, LX , but in clusters with very diﬀerent masses. We do however, see some diﬀerences in the overall luminosities of XRBs produced in star clusters. Figure 2-8 shows that the X–ray luminosities of binaries produced in clusters extend to higher values in the Antennae than in NGC 4449.

In this work, we have studied XRBs with secure parent cluster ages to investigate the production of high and intermediate-mass XRBs. We found tentative evidence that HMXBs and IMXBs form in diﬀerent ways. In an upcoming work, we present a new, complementary technique to separate HMXBs from IMXBs, one which uses the optical emission from the donor star in the XRB to estimate its mass. We will apply this method to the XRB population in the spiral galaxy M101, and see if the resulting X–ray luminosity functions of high and intermediate-mass XRBs also suggest that these populations formed through diﬀerent mechanisms.

In the cases where XRBs are embedded within a cluster, the age of the parent cluster provides strong evidence for they XRB classiﬁcation. Statistically, parent clusters with ages <10 Myrs are more likely to host HMXBs, since the most massive stars are the most dynamically active in the centers of dense star clusters (e.g., Garofali et al. 2012). If the XRB is found within a globular cluster (age >400 Myrs) it’s

36 most certainly a LMXB, since only low mass stars are left in these ancient stellar systems. Amazingly, galaxies that are within a distance of 10 Mpc are close enough where HST images can detect donor stars with masses down to ∼ 3M . In this way, we can measure the donor star masses directly, classify XRBs on an individual basis, and produce the most accurate XRB catalogues to date.

37 Chapter 3

New Methods for Separating High and Low Mass X–ray Binaries in

Star–Forming Galaxies with an

Application to M101∗

The high energy emission from nearby, star-forming galaxies is dominated by

X–ray binaries, where a neutron star or black hole is accreting mass from either a low mass (<∼ 3 M ) or high mass (>∼ 8 M ) star. Donor stars with intermediate masses ≈ 3 − 7 M tend to be more rare. It is not possible to separate low and high mass X–ray binaries (LMXBs and HMXBs) from X–ray observations alone. We have developed two techniques to distinguish high- from low-mass X–ray binaries, by comparing high resolution X–ray and optical images from Chandra and Hubble, respectively. HST observations of M101 are deep enough to directly detect donor stars down to ≈ 3 M ; therefore if no donor is detected, the source is a LMXB. A comparison between measured colors and luminosities of a detected donor star and

∗This chapter is a manuscript in preparation for submission to the Astrophysical Journal

38 theoretical isochrones distinguishes high-mass (M >∼ 8 M ) from intermediate-mass

(M = 3 − 8 M ) donor stars. We ﬁnd that ≈ 10% of the X–ray binaries in M101 still live within their parent

star clusters. In these cases, the age of the parent cluster provides strong evidence

for the type of binary. We have classiﬁed, on a source-by-source basis, X–ray point

sources in the nearby spiral galaxy M101, and ﬁnd that LMXBs have likely been a

major source of contamination in previous catalogs of HMXBs. We present the ﬁrst

‘clean’ luminosity functions for HMXBs, IMXBs, and LMXBs in a late-type galaxy,

and compare with previous results.

3.1 Introduction

In early-type, non star-forming galaxies, it is well established that the number of

low-mass X–ray binaries (LMXBs), where a neutron star or black hole is accreting

from a low-mass (<∼ 2−3 M ) donor star, correlates with the total stellar mass of the galaxy (Gilfanov 2004). We know signiﬁcantly less about X–ray source populations in late-type, star-forming galaxies, such as the frequency and spatial distribution of LMXBs, largely due to their complicated star formation histories and dust-ridden morphologies. Because they have formed stars recently, late-type galaxies also include a population of short-lived high-mass X–ray binaries (HMXBs), where the donor is a massive star (>∼ 8 M ). The work of Mineo et al. (2012) makes an early attempt at broadly classifying

XRBs in late–type galaxies based on their location within the galaxy. For example,

XRBs found in the bulge region of M101 are classified as LMXBs while XRBs found in the disk region are classified as HMXBs. The more sophisticated method outlined in Lehmer et al. (2017) classifies XRBs by dividing the galaxy into regions with

39 diﬀering star–formation rates and histories, arguing that HMXBs (LMXBs) are more likely to be found in regions with high (low) star–formation rates.

The seminal work by Grimm et al. (2003) established that the HMXB XLF has a near-universal shape, and that its normalization broadly scales with the global star formation rate of the host galaxy. Figure 3-1, from state-of-the-art work by Mineo et al (2012), shows that the cumulative X–ray luminosity functions of 29 nearby late- type galaxies all have fairly similar shapes, which could be described by a Schechter function, ψ(L) ∝ Lαe−L/L∗ , with an index α ≈ −1.6 and a marginally signiﬁcant

40 −1 break near LX ∼ 10 erg s (the solid line is an imposed broken power-law), and no evidence for a break at the Eddington limit for a stellar mass compact object

(neutron star or black hole). Somewhat surprisingly, the larger sample considered by

Mineo et al. (with respect to the original work by Grimm) did not yield the expected reduction in the relation scatter, which remains signiﬁcant at σ ≈ 0.4 dex. Mineo et al. (2012) and others argue that the scatter has a physical origin, likely variations in the star formation histories between galaxies or metallicity-dependent formation of

HMXBs.

Our understanding of LMXB populations comes from early-type galaxies. In early-type galaxies, ≈ 20-70% of LMXBs are found in ancient globular clusters rather than in the ﬁeld (e.g., Angelini et al. 2001; Jordan et al. 2004; Kundu et al. 2002,

2007; Sarazin et al. 2003; Humphrey & Buote 2008). Also, the shape of the XLF diﬀers from that of HMXBs, and can be described by a Schechter function with a

38 power-law index of ≈ −2 and downturn at a luminosity LX ≈ few × 10 erg/s (Kim & Fabbiano 2004).

One major, under-appreciated deﬁciency in all previous works has been the in- ability to directly separate HMXBs from LMXBs in star-forming galaxies. The well-known LMXB-galaxy mass scaling relation determined from early-type galaxies doesn’t work in late-type galaxies (Mineo et al. 2012). Figure 1-3 visually demon-

40 Figure 3-1: The cumulative X–ray luminosity functions for 29 nearby star-forming galaxies, normalized by their star– formation rates. The solid line shows the best-ﬁt broken power-law scaling relation. The scatter in the relation remains fairly high, ∼ 0.4 dex. Credit: Mineo et al. 2012

41 strates the assumptions and methodology used by Mineo et al. (2012), and their limitations. They presented the speciﬁc case of M101. X–ray sources within the inner circular region and outside of the outer circular region (corresponding to the

D25 radius of the galaxy) were excluded from their analysis, as the former (red circles) were assumed to be predominantly LMXBs associated with the bulge, while the latter (dark blue squares) were believed to be dominated by background sources

(mostly AGN). The resulting luminosity functions are entirely based on the X–ray sources identiﬁed by the small cyan crosses located between the inner and outer circles, which are assumed to be HMXBs in the disk. Interestingly, although Mineo et al. attempted to correct for LMXB contamination, they found that the LMXB

XLF established by Gilfanov (2004) for early-type galaxies signiﬁcantly over-predicted the contribution of LMXBs, and that its subtraction would result in negative values

(see section 7.3 in Mineo et al. 2012). This suggests that the LMXB populations in early- and late- type galaxies are characterized by diﬀerent XLFs, either in shape or normalization.

We develop direct, new techniques to separate HMXBs and LMXBs on a source- by-source basis. Many images of nearby galaxies (<∼ 10 Mpc) taken in the optical with the Hubble Space Telescope are deep enough to directly detect donor stars down to

M ≈ 3 M . This means that if no donor star is observed for an XRB, it is a LMXB. If a donor star is detected, its colors and luminosities can be compared with model predictions to estimate it’s mass, and sources where the donor has M >∼ 8 M are

HMXBs. Those with donors between 3 − 8 M are intermediate mass XRBs. Note that the LMXBs selected via our criteria may diﬀer somewhat from those in early type galaxies, where the donor star masses are somewhat lower, M <∼ 1 − 2 M . Nonetheless, we can take the next major step in understanding XRB populations in star-forming galaxies by applying our method. Some XRBs still reside in their parent clusters. In these cases, the age of the cluster gives information on the type of binary:

42 clusters younger than ≈ 50 Myr likely have a high mass donor (which are the most

dynamically active in dense star clusters) making them a HMXBs, and those older

than ≈ 400 Myr have donors with masses less than 3 M , making them a LMXB. Here, we present our methods and apply them to the X–ray point sources in M101.

We assume a distance of 6.4 Mpc to M101 (Shappee & Stanek 2011), based on recent

Cepheid work. This gives a physical scale where 100 ≈ 31 pc.

The rest of this chapter is organized as follows. Section 2 describes the catalog of

X–ray point sources from Chandra ACIS observations of M101. Section 3 presents the optical HST observations, the selection of candidate donor stars and how we estimate their masses, as well as parent star clusters and how we estimate their ages. Here, we classify each X–ray source in M101. Section 4 presents the spatial distributions and general properties of the LMXBs, IMXBs, and HXMBs. Section 5 presents X– ray luminosity ﬁts to each of these populations, and Section 6 summarizes the main results of this work.

3.2 Catalog and X–ray Source Classiﬁcation

3.2.1 Catalog of X–ray Binaries in M101

The Chandra X–ray Center has recently released their second source catalog for observations through the end of 2014 (see Evans et al. 2010), providing a homogeneous catalog of thousands of X–ray sources in tens of nearby late-type galaxies. The second release source catalog is thus an excellent starting point for this project. We selected an initial X–ray source list by requiring a > 3σ detection, and excluding sources that

are saturated or extended, and ﬁnd that 140 X–ray sources fall within the ≈ 100 × 100

optical mosaic of M101 described in the next section. Table 3.1 compiles the RA,

DEC, and total X–ray ﬂux for each of these sources.

43 3.2.2 HST Observations and Alignment

M101 has been observed with HST through several diﬀerent programs. Here, we

use the 10 pointings observed in the B,V, and I bands with the Wide Field Channel

(WFC) of the Advanced Camera for Surveys (ACS) as part of program GO-9490 (PI:

K. Kuntz). Each pointing was observed with two exposures for a total of 15 minutes in the B band, and 12 minutes in each of the V and I bands. No dithering was

performed. We show a V band mosaic of the 10 ﬁelds in Figure 3-2. This covers

≈ 100 × 100 centered on the nucleus of M101. The white lines are the gaps between the two detectors that make up the ACS/WFC, where no data is available.

We used the following procedure to align the Chandra X–ray source coordinates and the optical mosaic. We identify AGN scattered throughout the HST image using the GAIA database via the interactive ESASky (http://sky.esa.int), and use this to shift the optical mosaic to the World Coordinate System (WCS). Next, we identify

13/14 background galaxies and foreground stars (X11, X39, X55, X66, X80, X101,

X113, X116, X114, X171, X166, X177, and X180), and ﬁnd that the X–ray sources are shifted by 0.600 from their optical counterparts, and shift the coordinates to align the optical and x-ray coordinate systems. We ﬁnd that the random errors are fairly small, ≈ 0.300, from the standard deviation of the separations between the X–ray and optical coordinates for these 13/14 sources after shifting. Our estimate of the astrometric uncertainty in the X–ray coordinates in M101 is similar to that found previously for the Antennae by Rangelov et al. (2013).

In Figure 3-3, we show a 1.700 ×1.700 portion of an optical BVI color image centered

at the location of each X–ray source. Positional uncertainties of 1σ = 0.300 and

2σ = 0.600 are shown as the two centered circles. Potential donor stars or parent

clusters are shown as the smaller circles. We ﬁnd that the thumbnails can be broadly

characterized in one of three ways. Some X–ray sources do not have any obvious

44 F5 F6 F4

F7 F1

F2 F10

Figure 3-2: HST V–band mosaic of the 10 ﬁelds in M101.

45 optical counterpart detected within the 2σ positional uncertainty (e.g., X1, X3, and

X5). In others, a clear, dominant single source is detected within 2σ (e.g., X18, X22,

and X32), while the rest have multiple optical sources within 2σ (e.g., X7, X12, and

X24). When there is no detected optical counterpart, it is very likely that there is a

low-mass donor that is below the detection limit, and hence the source is a LMXB.

We use the properties of the optical counterparts within 2σ to classify each X–ray source. Optical counterparts to X–ray sources will be one of the following types of objects: foreground star, donor star in M101 (X–ray binary), compact star cluster host (X–ray binary), supernova remnant, or a quasar or AGN in a background galaxy.

Quasars and AGN are are relatively easy to select, because they are either extremely red (far redder than expected of a red supergiant) point sources, or clearly from a background galaxy. Typically, when there are multiple potential optical counterparts, it’s because the X–ray source is in a region of active star formation, with lots of young massive stars, and the source is likely a HMXB. Star clusters are broader than the point-spread-function, and are distinguished from an individual or close pair of stars from their proﬁles (see description in Chandar et al. 2010, for example). The remaining counterparts are the donor stars in the X–ray binary.

3.3 Methods for Classifying X–ray Binaries

3.3.1 Classiﬁcation Based on Donor Star Masses

For each X–ray source, we identify all optical sources located within 2σ, and perform aperture photometry to estimate their luminosities in the available ﬁlters

(see the technical description in e.g., Chandar et al. 2010). We perform aperture photometry on each detected optical source that (has a center which) lies within the

2σ circle, using a 3 pixel aperture. We correct for missing ﬂux by using the encircled

46 X1L X2L X3L X4S X5L

X7I X8L X9L X10L X11I

X12H X13L X14L X15L X16I

X18S X19L X20H X21I X22H

X23L X24H X25H X26A X27L

X28L X29L X30H X31I X32I

X33I X34I X35H X36I X37I

Figure 3-3(a):

47 X38L X39H X40L X41L X42H

X43I X44H X45L X46H X47L

X48I X49H X50I X51H X52L

X53I X54I X55H X56H X57L

X58A X59H X60H X61H X62I

X63H X64H X65L X66H X67I

X68H X69L X70H X71H X72H

Figure 3-3(b):

48 X73H X74H X75H X76H X77I

X78L X79L X80A X81L X82L

X83A X84L X85L X86I X87L

X88L X89L X90L X91A X92L

X93H X95A X96I X97L X98H

X99L X100H X101L X102A X104H

X105H X106L X108L X109I X110I

Figure 3-3(c):

49 X111I X112S X113L X114I X115H

X116I X117H X118I X119A X120H

X121I X122I X123I X124Q X125H

X126H X127A X128I X129H X132L

X134S X140L X141L X143S X145L

X146I X149I X150H X151H X155I

X156H X158L X160S X161L X162H

Figure 3-3(d):

50 energy distribution tables for point sources given in Sirianni et al. (2005), and convert

to the VEGAMAG system by using the zeropoint for each ﬁlter that is given on

the ACS instrument page (http://www.stsci.edu/hst/acs/analysis/zeropoints). We

record the V magnitude and the V − I color for each source in Table 3.1 (a number of sources are not detected in the B band, since it doesn’t go as deep).

We ﬁnd that we can detect optical sources down to mV ≈ 28 mag, which is deep enough to detect a ≈ 3 M star at the distance of M101. This is a critical point: if we cannot detect a donor star directly in the optical image, then the source is very likely to be a LMXB. We have checked the HST observations of M101, and this is true throughout the images, except in a few small portions in the most crowded part of the spiral arm and nuclear regions. In M101, we ﬁnd that there is no detectable donor star in 47 of the Chandra detected XRBs.

When a point source is detected, we compare the colors and magnitudes with theoretical isochrones in order to estimate the mass of each donor. In Figure 3-4, we compare the measured colors and luminosities of potential donor stars within

00 2σ = 0.6 with theoretical mass tracks of 1, 2, 3, 5, 8, 20, and 40 M from the Padova models (Marigo et al. 2008, Girardi et al. 2008, Bertelli et al. 1994) at solar metallicity (available here: http://stev.oapd.inaf.it/cmd).

Solar metallicity is appropriate for the radii covered by our study, since the well- known metallicity drop in M101 occurs a bit further out. We correct the photometry for foreground extinction, which is quite low towards M101, with AV ≈ 0.02 mag and E(V − I) ≈ 0.01 mag. This ﬁgure shows several interesting results. First, the

lowest luminosity donor stars that we detect reach down to ≈ 3 M , and the masses

extend up from there, to ≈ 40 M . Most detected donor stars have masses higher

than >∼ 8 M , but there are a few between 3 − 7 M . If these do not suﬀer from a signiﬁcant amount of extinction, then they are candidate intermediate-mass donors,

but if they do, then they could end up having estimated masses >∼ 8 M . We examine

51 Figure 3-4: Potential donor stars in M101 compared to theoretical mass tracks from the Padova models at solar metallicity.

52 the BVI color images of each potential intermediate-mass donor, and ﬁnd that there

is very little visual evidence for dust lanes or dark patches that might cause extinction,

and which show up nicely in other parts of the mosaic, with the exception of X11,

X34, X110, X111, X122, X123, and X128.

The classiﬁcation for each source is given in Table 3.1. When no optical source is

detected within 2σ = 0.600, the source is classiﬁed as a LMXB. Most of these sources

appear to be in regions of M101 with low background and extinction. Sources with

a donor star within 2σ that has a mass consistent with the 8 M or higher isochrone are classiﬁed as a HMXB. Note that there are a number of XRBs in crowded regions,

where it isn’t possible to tell exactly which star is the donor. In all but a few of these

cases, all potential donor stars have estimated masses of at least 8 M , so we classify them as a HMXB.

We also assign each X–ray source a ‘confidence’ flag, where 1 =high confidence,

2 =medium conﬁdence, and 3 =low conﬁdence (Table 4.1). X–ray sources with one

candidate donor within 2σ in a region of low extinction will be assigned a Flag=1.

LMXBs with a high background are given lower conﬁdences (Flag=2) since there

is a possibility of an intermediate mass star going undetected. Similarly, IMXBs

detected in regions with high backgrounds are given lower conﬁdences (Flag=2) since

the photometry is less certain. LMXBs are given a low conﬁdence (Flag=3) if a dim

optical source is detected exactly 2σ from the X–ray source. HMXBs and IMXBs

are given low conﬁdences (Flag=3) if there are multiple optical sources within 2σ, we cannot identify the donor, and all optical sources have vastly diﬀerent mass ranges.

HMXBs and IMXBs are assigned a Flag=2 if there are multiple optical sources within

2σ, we cannot identify the donor, yet all optical sources fall within the same mass range. If the X–ray source is located within the nucleus of the galaxy, we assign it a low conﬁdence (Flag=3).

53 In this work, we do not make any explicit correction to for extinction to the

photometry. For our purposes, correcting for extinction would increase the luminosity

of some “intermediate-mass” 3 − 7 M stars enough to push them above 8 M , i.e. into the regime of HMXBs. From the optical images we assess, on a case-by-case basis, the level of extinction that we think may be aﬀecting the source, due to the presence of dark patches and dust lanes. For the vast majority of our sources, we do not see visual evidence of extinction.

54 Table 3.1: M101 X–ray Point Source Properties

X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

1 210.80459478 54.34798659 35.52±36.09 – – – LMXB 1 2 210.81203354 54.35288294 36.84±37.09 21.403± 0.018 0.763± 0.066 1.141± 0.037 LMXB 1 Old Globular Cluster 3 210.78921548 54.34797275 36.45±36.53 – – – LMXB 2 4 210.79067377 54.35415710 37.22±37.18 – – – Foreground Star 1 5 210.81177559 54.35640599 36.54±36.80 – – – LMXB 1 7 210.79470128 54.35730647 36.50±36.70 23.788± 0.067 -0.143± 0.173 -0.184± 0.238 IMXB 3 8 210.81871469 54.35156777 37.02±37.03 – – – LMXB 1 9 210.80807936 54.36014507 36.69±36.86 – – – LMXB 1 10 210.81296842 54.33833502 36.65±36.72 – – – LMXB 1

55 11 210.80680496 54.33595766 36.95±36.86 25.673± 0.117 0.579± 0.350 0.142± 0.304 IMXB 1 12 210.78339349 54.35642577 36.33±36.35 23.198± 0.034 0.020± 0.083 0.166± 0.089 HMXB 2 13 210.80707693 54.36301871 37.38±37.14 – – – LMXB 1 14 210.81095628 54.36444162 36.07±36.33 – – – LMXB 1 15 210.80544153 54.36585707 36.71±37.01 – – – LMXB 1 16 210.77579911 54.33658448 37.15±37.24 26.675± 0.149 **** 2.287± 0.219 IMXB 1 18 210.83026118 54.36368372 37.13±36.90 – – – Foreground Star 1 19 210.82291324 54.36843285 37.08±37.12 – – – LMXB 1 20 210.77915586 54.36927218 37.13±37.06 26.042± 0.132 0.359± 0.389 0.409± 0.356 HMXB 1 21 210.84444158 54.34407150 36.42±36.63 22.941± 0.040 0.441± 0.129 1.015± 0.083 IMXB 3 few 100 Myr cluster 22 210.78086832 54.32721238 38.04±37.61 24.596± 0.058 0.686± 0.194 1.040± 0.113 HMXB 1 23 210.76094602 54.35570284 36.94±37.18 – – – LMXB 1 24 210.77502722 54.32913878 37.00±36.94 23.447± 0.043 -0.054± 0.100 -0.039± 0.120 HMXB 2 25 210.75549539 54.35924087 36.57±36.82 24.806± 0.096 0.866± 0.369 1.708± 0.145 HMXB 2 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 26 210.77750088 54.37454735 37.17±37.18 25.585± 0.138 0.432± 0.402 2.360± 0.191 AGN 1 27 210.80711419 54.31912856 37.20±36.88 – – – LMXB 1 28 210.75331483 54.33813272 36.16±36.41 – – – LMXB 1 29 210.85490106 54.34717202 37.20±36.97 – – – LMXB 1 30 210.76975287 54.32349042 35.34±35.94 20.287± 0.010 0.147± 0.027 0.379± 0.025 HMXB 2 31 210.75599698 54.36568047 35.23±35.84 26.508± 0.142 1.914± 0.837 1.660± 0.242 IMXB 1 32 210.74769287 54.34518583 36.70±36.98 26.391± 0.132 -0.629± 0.311 1.038± 0.341 IMXB 1 33 210.75209971 54.33396490 36.35±36.57 24.743± 0.069 -0.074± 0.175 0.294± 0.207 IMXB 1 34 210.85553408 54.33771635 37.24±37.17 26.085± 0.142 0.602± 0.454 1.343± 0.249 IMXB 1

35 210.85768551 54.35690851 36.00±36.35 25.035± 0.103 0.287± 0.257 0.914± 0.231 HMXB 1 On 8 M line 36 210.86109840 54.34518277 36.74±36.89 25.843± 0.117 0.224± 0.373 0.653± 0.314 IMXB 3 Cluster

56 37 210.78774492 54.38302663 37.08±37.08 25.732± 0.126 1.060± 0.591 2.120± 0.192 IMXB 1 extended in I 38 210.83816797 54.31888386 36.96±37.30 – – – LMXB 2 39 210.85909281 54.33075178 36.79±37.16 22.357± 0.021 -0.061± 0.053 0.011± 0.063 HMXB 2 40 210.84008322 54.37935524 36.82±37.09 – – – LMXB 1 41 210.87024789 54.34968002 36.42±36.76 – – – LMXB 1 42 210.86986355 54.35357395 36.51±36.68 21.556± 0.016 0.194± 0.043 0.330± 0.041 HMXB 1 blended with faint star 43 210.86625431 54.33456839 36.71±36.88 23.984± 0.062 0.349± 0.197 1.016± 0.134 IMXB 1 few 100Myr cluster 44 210.78278101 54.31014354 37.23±37.10 25.032± 0.080 0.811± 0.268 1.112± 0.159 HMXB 1 45 210.81980715 54.30967130 36.83±36.91 – – – LMXB 2 46 210.82389077 54.30995830 37.69±37.28 25.723± 0.134 -0.045± 0.349 1.222± 0.275 HMXB 2 47 210.86117856 54.37245444 36.59±36.84 – – – LMXB 1 48 210.77202725 54.38653937 36.64±36.75 27.753± 0.159 0.460± 0.927 3.670± 0.209 IMXB 1 49 210.73182869 54.34130163 36.97±37.48 23.666± 0.048 -0.095± 0.114 0.031± 0.139 HMXB 1 50 210.87467999 54.34931157 36.52±36.73 25.510± 0.100 0.074± 0.257 0.021± 0.372 IMXB 1 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 51 210.87078105 54.36364921 36.23±36.57 26.216± 0.136 -0.133± 0.376 -2.331± 0.545 HMXB 1 52 210.86926347 54.36719895 36.48±36.90 – – – LMXB 1 53 210.73475045 54.32977803 36.65±36.82 25.698± 0.113 0.180± 0.273 0.477± 0.272 IMXB 1 54 210.81683350 54.30540505 37.14±37.13 24.259± 0.075 0.390± 0.227 0.535± 0.191 IMXB 3 few 100Myr cluster 55 210.81385375 54.39361678 36.78±37.10 **** **** **** HMXB 1 Supergiant. No V-band photometry 56 210.78601281 54.39344450 35.84±36.16 23.966± 0.046 -0.179± 0.110 -0.285± 0.154 HMXB 2 57 210.80821596 54.30282152 37.27±37.65 – – – LMXB 2 58 210.78487416 54.30319033 37.25±37.21 25.380± 0.107 0.759± 0.438 2.414± 0.153 AGN 1 59 210.80961551 54.30171040 38.31±38.73 23.258± 0.041 -0.037± 0.098 0.169± 0.113 HMXB 1 60 210.75875767 54.39023855 36.19±36.51 24.354± 0.055 0.306± 0.179 0.515± 0.157 HMXB 2 61 210.87486082 54.37464663 37.31±37.52 22.905± 0.037 0.066± 0.099 -0.239± 0.116 HMXB 2

57 62 210.84573878 54.30612660 36.04±36.41 24.686± 0.077 0.255± 0.228 0.905± 0.188 IMXB 1 few 100Myr cluster 63 210.80269914 54.29896462 36.01±36.31 22.688± 0.030 -0.174± 0.070 -0.256± 0.094 HMXB 2 64 210.72877795 54.37429839 36.89±37.31 23.081± 0.030 -0.133± 0.074 -0.185± 0.096 HMXB 2 65 210.87806505 54.37267451 37.63±37.28 – – – LMXB 1 66 210.80193387 54.29825264 37.15±36.86 22.803± 0.034 0.106± 0.082 0.264± 0.093 HMXB 2 67 210.71509138 54.34796645 36.37±36.62 25.682± 0.111 0.841± 0.421 2.118± 0.162 IMXB 1 68 210.86581629 54.31305996 36.71±37.00 23.862± 0.053 -0.113± 0.127 0.080± 0.151 HMXB 1 69 210.79379239 54.29703434 36.80±36.78 – – – LMXB 1 70 210.71280908 54.34491676 36.30±36.62 24.213± 0.054 0.108± 0.143 -0.060± 0.165 HMXB 1 71 210.85014185 54.39360325 36.80±36.94 24.018± 0.059 -0.110± 0.134 -0.285± 0.184 HMXB 1 72 210.86323647 54.30868588 37.30±37.14 21.063± 0.013 1.161± 0.053 1.049± 0.026 HMXB 1 Red Supergiant 73 210.73980026 54.30943856 36.89±37.03 24.680± 0.067 -0.334± 0.152 -0.241± 0.210 HMXB 1 74 210.75486796 54.39484672 36.79±36.94 23.904± 0.047 0.138± 0.129 0.269± 0.127 HMXB 1 75 210.88708755 54.37043242 35.94±36.56 24.838± 0.097 -0.281± 0.200 0.013± 0.264 HMXB 2 Frame cut oﬀ X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 76 210.77176774 54.29781155 37.55±37.37 23.299± 0.037 0.357± 0.104 1.117± 0.076 HMXB 1 77 210.87770210 54.38045196 36.64±36.90 25.520± 0.128 0.049± 0.283 0.359± 0.330 IMXB 1 78 210.72580554 54.31723309 37.29±37.33 – – – LMXB 2 79 210.71582026 54.32569266 36.57±36.99 – – – LMXB 2 80 210.89711735 54.35657179 36.92±37.09 26.757± 0.220 1.138± 0.873 2.541± 0.305 AGN 1 extended in I 81 210.85526954 54.39523338 37.31±37.28 – – – LMXB 2 82 210.80222983 54.40499822 37.17±37.02 – – – LMXB 1 83 210.72226410 54.31554729 36.33±36.50 – – – AGN 1 84 210.70688228 54.33323370 36.82±36.83 – – – LMXB 1 85 210.71410311 54.37542647 36.38±36.73 – – – LMXB 1 86 210.71369316 54.32171265 36.74±37.03 25.594± 0.116 -0.012± 0.306 0.360± 0.335 IMXB 1 two blended stars

58 87 210.79642295 54.29039124 36.51±36.78 – – – LMXB 2 88 210.87038673 54.39228904 35.85±36.48 – – – LMXB 2 89 210.76405312 54.29357106 36.70±36.83 – – – LMXB 2 90 210.70108101 54.34093576 36.62±36.81 – – – LMXB 1 91 210.89962087 54.36738890 36.41±36.82 – – – AGN 1 92 210.73032979 54.39161623 36.23±36.64 – – – LMXB 2 93 210.82776933 54.40821670 37.74±37.32 24.732± 0.065 0.813± 0.238 1.966± 0.103 HMXB 1 Red Supergiant 95 210.73913038 54.29963523 36.80±37.09 25.119± 0.101 3.042± 1.286 2.863± 0.128 AGN 3 96 210.90759382 54.34406218 36.79±37.06 26.413± 0.206 0.289± 0.559 1.231± 0.421 IMXB 3 few 100Myr cluster 97 210.82925794 54.28876685 36.51±37.09 – – – LMXB 2 98 210.90012735 54.32354306 38.17±37.85 25.025± 0.084 -0.178± 0.199 -0.145± 0.234 HMXB 1 99 210.88962959 54.38531463 36.91±37.21 – – – LMXB 2 Red source only in I band 100 210.73913207 54.29801390 36.60±36.78 23.046± 0.032 -0.135± 0.081 -0.038± 0.101 HMXB 1 101 210.77981649 54.41015113 35.71±36.16 – – – LMXB 1 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 102 210.86553267 54.40002228 37.26±37.19 – – – AGN 1 104 210.81864085 54.28550109 36.13±36.55 24.585± 0.073 -0.200± 0.166 -0.112± 0.213 HMXB 1 105 210.81456628 54.28443102 36.97±37.12 22.358± 0.033 0.797± 0.131 -0.505± 0.135 HMXB 2 Young 3–5Myr cluster 106 210.71001488 54.38507362 37.12±37.14 – – – LMXB 2 108 210.74914880 54.40709003 36.48±36.80 – – – LMXB 2 109 210.70463775 54.31164367 37.01±36.92 20.502± 0.011 0.057± 0.029 0.593± 0.025 IMXB 3 few 10s of Myr compact cluster 110 210.83945839 54.28409742 36.85±37.12 25.019± 0.078 -0.121± 0.201 0.271± 0.241 IMXB 3 111 210.86833875 54.40512702 35.84±36.49 26.675± 0.148 -0.685± 0.339 -0.448± 0.462 IMXB 1 112 210.78502515 54.41639559 37.71±37.49 – – – Foreground Star 1 113 210.70296924 54.38642848 36.62±36.90 – – – LMXB 2 114 210.68400126 54.34297283 36.50±36.82 25.877± 0.105 -0.533± 0.232 1.061± 0.274 IMXB 3

59 115 210.87895321 54.29539378 35.00±35.54 23.529± 0.036 0.213± 0.115 1.113± 0.075 HMXB 2 116 210.88890563 54.29992144 37.05±37.18 26.160± 0.138 -0.333± 0.313 0.804± 0.406 IMXB 3 24.743± 0.100 -0.213± 0.214 0.120± 0.279 Multiple 117 210.68270016 54.33491545 36.84±37.00 24.226± 0.062 -0.064± 0.151 -0.254± 0.190HMXB 3 Potential 25.452± 0.105 -0.185± 0.245 -0.014± 0.286 Donors 118 210.70125190 54.30798340 36.78±36.90 25.432± 0.092 0.305± 0.278 0.576± 0.211 IMXB 1 119 210.90106592 54.30408902 37.81±37.25 – – – AGN 1 Elliptical Galaxy 120 210.88443088 54.40386885 36.43±36.85 24.864± 0.066 0.638± 0.215 0.891± 0.139 HMXB 1 121 210.71414651 54.29676623 37.32±37.15 25.309± 0.101 -0.023± 0.234 0.374± 0.246 IMXB 3 122 210.67816566 54.33781145 36.66±36.91 25.348± 0.112 0.546± 0.292 0.885± 0.229 IMXB 1 123 210.86123572 54.28215899 36.90±37.21 26.941± 0.169 0.195± 0.433 0.191± 0.366 IMXB 1 124 210.83543333 54.27565217 37.73±36.99 24.950± 0.083 -0.061± 0.205 -0.524± 0.283 Quasar Candidate 3 Red source only in I-band 125 210.89894820 54.40025523 36.58±36.48 25.252± 0.088 0.231± 0.231 1.156± 0.171 HMXB 2 126 210.77128173 54.42285314 36.69±36.80 25.823± 0.126 **** 2.814± 0.154 HMXB 1 Red Supergiant X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 127 210.91389098 54.30764096 37.26±37.15 – – – AGN 1 128 210.92470390 54.31924715 36.35±36.59 26.657± 0.145 2.362± 0.837 1.740± 0.249 IMXB 3 129 210.79072505 54.27101538 36.82±37.05 24.612± 0.070 0.078± 0.189 0.129± 0.218 HMXB 1 blended stars 132 210.94204046 54.36599325 35.77±36.13 – – – LMXB 2 Red source only in I–band 134 210.89811721 54.28573529 38.12±37.85 – – – Foreground Star 1 140 210.95110139 54.35478793 36.38±36.58 – – – LMXB 1 141 210.81724163 54.43568296 37.14±37.37 – – – LMXB 1 143 210.94745104 54.37488239 37.11±37.12 – – – Foreground Star 1 145 210.76982219 54.25950349 37.11±37.26 – – – LMXB 2 146 210.91189135 54.28195928 37.33±37.21 26.966± 0.196 1.558± 0.874 1.436± 0.353 IMXB 1 149 210.95392701 54.37852061 36.24±36.48 26.878± 0.300 0.768± 1.007 2.075± 0.368 IMXB 1 hot pixel visible

60 150 210.93653091 54.40040714 37.67±37.33 25.007± 0.091 0.827± 0.291 1.855± 0.133 HMXB 1 hot pixel visible 151 210.96371322 54.34743304 37.13±37.16 23.083± 0.028 0.109± 0.080 0.207± 0.083 HMXB 1 155 210.66204185 54.29893782 37.42±37.40 25.387± 0.117 0.269± 0.371 0.754± 0.269 IMXB 3 22.839± 0.040 0.104± 0.085 -0.295± 0.104 Multiple 24.215± 0.083 1.098± 0.388 1.791± 0.108 Potential 156 210.96636761 54.36376970 37.16±37.32 HMXB 2 23.351± 0.045 -0.188± 0.088 0.172± 0.107 Donors 24.891± 0.114 -0.328± 0.190 -0.096± 0.237 158 210.90295808 54.27075074 36.62±36.87 – – – LMXB 1 160 210.63546753 54.33369052 38.04±37.69 – – – Foreground Star 1 161 210.94461876 54.40242372 36.53±36.83 – – – LMXB 1 162 210.96941626 54.36976485 36.25±36.49 29.153± 0.446 **** 5.794± 0.457 HMXB 1 3.3.2 Classiﬁcation Based on Parent Clusters Ages

We ﬁnd 9 (∼ 10%) of the X–ray sources in M101 have an optical counterpart that

is a compact star cluster rather than an individual donor star. Aperture photometry

for the clusters was performed in the same way as for the donor stars, in a 3 pixel

radius aperture. The only diﬀerence is in the aperture corrections, which depend on

the size of the cluster. We follow the methodology of Cook et al. (2018) to map the

concentration index, i.e. the diﬀerence in magnitude measured in 3 and 0.5 pixels,

onto an aperture correction.

At the distance of M101, clusters are too crowded to identify the donor star.

However, it is possible to use the age of the cluster to separate HMXBs from LMXBs.

In Figure 3-5, we compare the measured B − V and V − I colors of the clusters with predictions from the solar metallicity stellar evolutionary models of Bruzual &

Charlot (2003). This ﬁgure shows the evolution of cluster colors over time, as the most massive clusters stars ﬁnish their lives. The cluster colors have been corrected for a foreground reddening of E(B − V ) = 0.02, assuming a standard Milky Way-

type extinction law (Fitzpatrick 1999). The solid (dashed) line maps the predicted

evolution of a cluster with solar metallicity (1/5 solar metallicity). The solid circles

represent the 9 clusters hosting XRBs in M101. The star symbols show the predicted

locations for clusters at diﬀerent ages.

High–mass (>8 M ) and intermediate–mass (3–8 M ) stars have hydrogen burning lifetimes of ≈ 10 Myr and ≈ 300–400 Myr respectively. Therefore, XRBs found

in clusters >400 Myr must therefore be LMXBs since only low–mass stars at this age

will remain. N–body simulations show that clusters <10 Myr are too young to form

binary systems through dynamical interactions (Garofali et al. 2012). This, coupled

with the fact that the ﬁrst binaries to form primordially are high-mass, XRBs found in

clusters <10 Myr are more likely to be HMXBs. Clusters 10–400 Myr are suﬃciently

61 -0.5

1 0.0 10 100

V 1000 - 0.5 B

1.0

1.5 -1.0 -0.5 0.0 0.5 1.0 1.5 2.0 V - I

Figure 3-5: A comparison of measured host cluster B −V and V −I colors in M101 with predictions from the solar metallicity stellar evolutionary models of Bruzual & Charlot (2003). The solid (dashed) line maps the predicted evolution of a cluster with solar metallicity (1/5 solar metallicity). The solid circles represent the 9 clusters hosting XRBs in M101. The star symbols show the predicted locations for clusters at diﬀerent ages.

62 old enough that they no longer host the most massive of stars (>8 M ). At this age, the more massive of the remaining stars have had suﬃcient time to form binary pairs through dynamical interactions (Garofali et al. 2012). We therefore classify any

XRBs found in clusters 10–400 Myr as IMXBs.

Nebular line emission is obvious in the morphology of one of the clusters, which means that it is younger than ≈ 6 − 7 Myr, and still has massive, young stars. We therefore designate this X–ray source (X105) as a HMXB. We estimate the age of each cluster by ﬁnding the closest model to its colors. In one case, the measured colors are quite red, similar to those of ancient globular clusters in the Milky Way. We designate this X–ray source (X2) as a ‘LMXB’, since the only potential donor stars in ancient globular clusters have low masses, <∼ 1 M . There are several clusters (X21, X36, X43, X54, X62, X96, and X109) which have colors which are intermediate between

those of ancient globular clusters and clusters younger than a few tens of Myr, and

these likely host IMXBs. The cluster age estimates and resulting classiﬁcation of

their XRBs are noted in the table. Our ﬁnal catalog contains 44 HMXBs, 33 IMXBs,

and 47 LMXBs.

3.4 Spatial Distributions of XRBs

In the HST mosaic image of M101 presented in Figure 3-6, we show the locations of HMXBs (blue), IMXBs (green), and LMXBs (red) in M101, as well as AGN

(magenta), based on our source-by-source classiﬁcation. It is interesting to compare these results with Figure 1-3. The same 2 circles are drawn, one denoting the “bulge” region and the other the D25 radius. Mineo et al. (2014) assumed that the “bulge” region contains only LMXBs associated with the old stellar population. Although we

ﬁnd that this region is dominated by LMXBs, IMXBs and HMXBs are also found in the “bulge” region. The region between the 2 circles was assumed to contain a pure

63 disk population of HMXBs. Here, we again ﬁnd a mix of low and high mass XRBs.

Finally, Mineo et al. (2014) assumed that the X–ray sources beyond the outer circle were highly likely to be background sources, typically AGN. We find that this region is actually dominated by XRBs associated with M101. One of the more interesting results of this work is that LMXBs are found throughout M101. If they were restricted to a bulge or halo distribution, we would expect their radial profiles to be centrally concentrated, and to drop off sharply with distance. Instead, we continue to find

LMXBs beyond the D25 radius. Our results suggest that M101 has formed a population of LMXBs in the disk. This means that scaling relations developed for LMXBs in early-type galaxies cannot be applied directly to star-forming galaxies. We know that M101 contains an old disk population of star clusters (Simanton et al. 2015).

3.5 X–ray Luminosity Functions

Figure 3-7 shows one of the main results of this work. This ﬁgure compares the

X–ray luminosity functions (XLF) from our individually classiﬁed HMXBs (blue),

IMXBs (green), and LMXBs (red), and ﬁnds notable diﬀerences. The XLF for

HMXBs extends to significantly higher fluxes and has a flatter distribution than those for LMXBs and IMXBs. The distributions for IMXBs and LMXBs appear to be quite similar. Below, we present fits and quantitative results for these distributions.

We use three diﬀerent methods to ﬁt the X–ray luminosity functions of XRBs.

Method 1 presents the binned luminosity function with equal size bins in log ﬂux.

This is the simplest way to visualize the data, and provides insight into the shape of the distribution. We ﬁt a simple power law, dN/dL ∝ Lα to the binned distributions

(e.g., Chandar et al. 2010). Method 2 ﬁts a truncated power law to the cumulative luminosity distribution using the code MSPECFIT (Rosowlowsky 2005). This method returns the power-law index α, as well as a statistic, N0, which assesses whether or

64 15’’

Figure 3-6: HST mosaic image showing the locations of HMXBs (blue), IMXBs (green), and LMXBs (red) in M101, as well as AGN (magenta), based on our source-by-source classiﬁcation.

65 Figure 3-7: Cumulative X–ray luminosity functions for classiﬁed HMXBs (blue), IMXBs (green), and LMXBs (red) in M101.

66 not there is a statistically signiﬁcant downturn at the bright end of the distribution.

N0 values with a statistical significance of 3 or higher indicate a cutoff at the bright end of the luminosity function. Method 3 performs maximum likelihood fits of the

α Schechter function, ψ(L) ∝ L exp(−L/L∗) to the observed luminosity functions, and determines the best-ﬁt values and 1-, 2-, and 3-σ conﬁdence intervals for α and

L∗ (Mok et al. 2018). This method has the advantages of not requiring binned data (which can hide weak features at the ends of a distribution) or cumulative distributions (where the data points are not independent of one another), and provides robust constraints on any cutoﬀ M∗. The top panels in Figure 3-8 show the results from each of the three methods described above when applied to the total X–ray luminosity function of all XRBs in

M101 (HMXBs + IMXBs + LMXBs). The overall XLF appears to be well-described by a single power-law with α ≈ −2 in the left and middle panels. There is no detection of a Schechter-like downturn in the middle panel, since N0 = 2.58 ± 1.58 (∼ 1.5σ) indicates no detection of a cutoﬀ, or in the right panel, where the 2 and 3 σ contours do not close around any particular value, but rather gives a lower limit for L∗.

67 68

Figure 3-8: Statistical ﬁts to total XLF of XRBs in M101. The 3 rows of panels in Figure 3-9 show the ﬁtting results for HMXBs (1st row),

IMXBs (2nd row), and LMXBs (3rd row). In these cases, we show sources with 1+2 confidence levels, because we find that the results do not change significantly when only sources with a confidence level of 1 or of 1+2 are used. The luminosity functions of HMXBs are all well fit by a single power-law with α = −1.43±0.17 from the binned method and α = −1.70 ± 0.10 from the cumulative one. There is no evidence for a downturn at the bright end. For IMXBs, the luminosity functions are clearly steeper, with α = −2.05 ± 0.46 (α = −2.27 ± 0.27) from the binned (cumulative) method, again with no evidence for a downturn. LMXBs have a similarly steep luminosity function, with α = −1.71 ± 0.26 (α = −2.07 ± 0.11). For LMXBs there is no evidence for a downturn at the bright end.

69 70

Figure 3-9: Statistical ﬁts to XLF of HMXBs (top), IMXBs (middle), and LMXBs (bottom) in M101. 3.6 Discussion and Conclusions

The seminal work by Grimm et al. (2003) on HMXBs in star-forming galaxies

found a power-law index α ≈ −1.6. LMXBs in early-type galaxies follow a power law with a somewhat steeper index, closer to ≈ −2 at lower luminosities. At log(LX ) >∼ 37.0 − 37.5 the differential slope gradually steepens, and has an abrupt cutoff near log(LX ) ∼ 39.0–39.5. The value of the cutoff is significantly lower, by an order of magnitude, than that for HMXBs (e.g., Gilfanov 2004).

The XLFs found here based on our source-by-source classiﬁcation of HMXBs

and LMXBs in M101 are, to ﬁrst order, similar to those published previously for

HMXBs in star-forming galaxies and LMXBs in early-type galaxies, suggesting that

our classiﬁcation methods work(!)

We ﬁnd that in M101 the XLFs of IMXBs and LMXBs are quite similar to each

other, but distinct from those of HMXBs (See Figure 3-9). This may point to diﬀerent

formation mechanisms for XRBs with donor stars <8 and ≥8 M . We don’t see a a Schechter-like downturn at the high luminosity end for any subpopulation. We need a larger sample of galaxies to improve our statistics to see if eventually we can see a downturn.

71 Chapter 4

Populations of X–ray Binaries in the Spiral Galaxy M83

4.1 Background

M83 is a face–on grand-design SAB(s)c spiral galaxy in the Southern Hemisphere, often called the ‘Southern Pinwheel’ (Figure 4-1). It is well known for the starburst in its nuclear region (Figure 4-2) and its well-deﬁned spiral arms. M83 is located at a distance of 4.61 Mpc (Saha et al. 2006), where 100 ∼ 22pc. M83 has a typical star

−1 formation rate for a large spiral galaxy, with 3.35 M yr (Lee et al. 2009) and a B–band luminosity of -19.8 mag (Pilyugin et al. 2012).

In comparison, M101 is a face–on ﬂocculent SAB(rs)cd spiral galaxy. It’s found at a distance of 6.4 Mpc (Shappee & Stanek 2011), where 100 ∼31pc. M101 has a

−1 star formation rate of 4.52 M yr (Lee et al. 2009) and a B–band luminosity of -21.1 mag (Leitherer et al. 2011).

72 Figure 4-1: A composite RGB image of M83 taken with the HST .

73 Figure 4-2: A composite RGB image of the nuclear region of M83 taken with the HST .

74 4.2 Observations

4.2.1 X–ray Observations with Chandra

Observations of M83 were taken with the Advanced CCD Imaging Spectrometer

(ACIS) instrument on the Chandra X–ray Observatory (CXO) in order to show both

hard and soft X–ray emission. The observations have integration times ranging from

9–150 ks, giving a total exposure time of 790 ks for the useful data obtained by K.

Long (Prop ID. 12620596), G. Rieke (Prop ID. 1600489), and A. Prestwich (Prop ID.

267005758).

The X–ray sources used in this work were culled from the publicly available cat-

alogue of Long et al. (2014). Sources were identiﬁed using CIAO’s Mexican–hat

wavelet source detection routine wavdetect (Freeman et al. 2002). Over 400 point sources were detected in the spiral arms and star forming regions of M83. Long et al. (2014) published the characteristics of these X–ray point sources including their location and X–ray luminosity (LX). A list of X–ray sources in M83 and their corresponding R.A., Decl., and LX is found in Table 4.1.

4.2.2 HST Observations

Observations of seven ﬁelds covering ≈ 162”×162” of M83 were observed with the Hubble Space Telescope(HST) with the WFC3/UVIS camera. Here we focus on

BVI and Hα observations. The approximate field coverage is shown projected onto a ground-based V-band image of M83 in Figure 4-3. The two white boxes show the archival fields from the Early Release Science Program (ID: 11360; PI: R. O’Connell) taken on 2009 August and 2010 March, while the black boxes show archival fields from a cycle 19 HST General Observer program (ID: 12513; PI: W. Blair) taken from from 2012 July through early 2012 September. In this work, we assume a distance of

75 4.61 Mpc to M83 (Saha et al. 2006) and a 1”∼22pc, which allows for accurate stellar

photometry in crowded ﬁelds.

Typically, point sources detected with HST have astrometric accuracies that are

better than ∼0.05”, while Chandra images have astrometric accuracies about an order

of magnitude higher. We estimated the positional uncertainty of the X–ray sources

in M83 by comparing the location of background AGN to their optical counterparts.

We identiﬁed 22 AGN in the HST images which we use to estimate a 1σ positional

uncertainty of ∼0.3”.

In Figure 4-4 we show BVI color images centered on each X–ray source. The

two concentric green circles represent 1 and 2σ positional uncertainties (0.3” and

0.6”). The smaller circles identify candidate donor stars and smaller dashed circles

identify candidate donor clusters. Optical counterparts to X–ray sources will either

be a foreground star, supernova remnant, donor star (i.e., an XRB), cluster (i.e.,

an XRB), or a quasar or AGN in a background galaxy. A catalogue of supernova

remnants have been identiﬁed by Long et al. (2014) and have been removed from our

catalogue. AGN and quasars are easily identiﬁed since they are extremely red or have

very identiﬁable structure (e.g., X402, X405, X412). Note that there are a number

of XRBs with multiple optical counterparts. In these cases, we make an attempt to

identify the most likely donor star or cluster. Where possible, the candidate donors

are identiﬁed in Figure 4-4 by a red circle.

4.3 X–ray Binary Donor Catalogue

With the exception of previously identiﬁed quasars, AGN, SNRs, and foreground

stars, we assume that each X–ray point source is an XRB in M83. The HST data

are suﬃciently deep that we should be able to detect donor stars down to ∼3 M , i.e. very close to the LMXB/HMXB separation limit. We follow the same deﬁni-

76 Figure 4-3: Locations of the seven M83 ﬁelds. The white boxes (Fields 1 and 2) are archival ﬁelds from Proposal 11360, while black boxes come from Cycle 19 Program 12513. Credit: Blair et al. (2014).

77 X043L X047L X048H X049L X050L

X051H X052L X054L X055L X056H

X058H X059H X062I X064L X066H

X068H X071H X072H X073H X074I

X075L X076L X077H X079I X080L

X082I X084L X085I X086I X087I

X088L X090H X091H X092I X094I

Figure 4-4(a):

78 X096H X098L X099H X101L X102I

X103L X109L X111I X112L X113I

X114H X115I X117I X118H X120H

X122L X123I X124H X125L X126H

X130L X132I X133L X135I X137L

X138L X139H X143H X144I X145I

X146L X147I X148I X150L X152H

Figure 4-4(b):

79 X153I X154I X157L X158H X160I

X161I X162H X163H X164H X165A

X167H X168L X170L X173I X175I

X176L X177I X178I X179H X180H

X182L X185L X187L X188I X189H

X190I X191I X192I X193I X194H

X196I X197H X198H X200H X201L

Figure 4-4(c):

80 X203I X204H X206L X208H X209H

X210?

X210H X213H X214H X215L X216H

X206L X213H

X217I X218H X219L X220H X222H

X223H X224H X225L X226H X227I

X228H X229H X230H X231H X232H

X234H X236L X237L X239H X240H

X241H X242L X244L X245H X246H

Figure 4-4(d):

81 X247L X248L X251H X252I X254I

X257I X258I X259L X260L X264I

X265H X266H X267I X269I X270L

X273I X274H X276L X277L X278H

X280L X281I X282L X283I X284H

X285H X286H X289H X290L X291L

X293H X294H X295L X296H X298H

Figure 4-4(e):

82 X299L X300L X301H X302L X303I

X304L X307L X308L X309H X312I

X314H X315L X317H X318H X321H

X322L X323L X327H X328I X329H

X331I X332I X333I X335L X337H

X343L X344L X345I X346I X347H

X349L X351H X358I X363I X366I

Figure 4-4(f):

83 X370I X371H X377I X378H X381I

X384I X393L X395L X396L X397I

X402A X405A X408L X409L X412A

Figure 4-4(g): tions for donor star masses as in the previous two chapters. X–ray sources with no detectable optical counterpart and hence no detectable donor star down to ≈ 3 M are classified as LMXBs. With the rare exception, X–ray sources with multiple potential optical counterparts within 2σ are located in star–formation regions with lots of young massive stars, and are therefore typically classified as HMXBs. XRBs found within compact star clusters are classified using the cluster’s estimated age, using the method developed in Chapter 2. Star clusters are identified and age dated using the methods and catalogue found in Chandar et al. (2010).

4.3.1 Classiﬁcation Based on Donor Star Masses

For each X–ray source, all optical sources within 2σ (0.6”∼13.404 pc) are identiﬁed using the IRAF task DAOFIND and found ∼900 objects in the V–band images, which includes individual stars, close blends of a few stars, star clusters, and background galaxies. We note that F555W was used for the V–band in Field 1 and that F547M

84 was used for the V–band in Fields 2–7. We conducted circular aperture photometry

with the IRAF task PHOT assuming a 3 pixel aperture radius and a background

annulus with an inner and outer radius of 5 and 8 pixels, respectively. We converted

the instrumental magnitudes to the VEGAMAG magnitude system by applying the

following zeropoints: F555W=25.80, F547M=24.72, F438W=24.98, F336W=23.46,

F814W=24.67, F657N=22.36, taken from the instrument website. Aperture corrected

V–band absolute magnitude, (V − I), and (B − V ) colors for candidate stars and

clusters are listed in Table 4.1.

Aperture corrections are determined from the mean magnitude diﬀerence between

apertures of 3 and 20 pixels for ∼150 isolated high signal-to-noise (S/N) stars and

clusters. Based on measurements for WFC3, an additional ∼0.9 mag is needed to

correct the photometry from 20 pixels to inﬁnity in each ﬁlter (see Instrument Science

Report WFC3 2016-03 for exact corrections used for each ﬁlter).

In M83 we ﬁnd 45 X–ray sources with no detectable donor star, which are classiﬁed

as LMXBs. We estimate the donor star mass by comparing candidate donor photom-

etry to the Padova group’s mass tracks (Marigo et al. 2008, Girardi et al. 2008,

Bertelli et al. 1994). In Figure 4-5, we compare the measured colors and absolute

V–magnitude of potential donor stars within 2σ with theoretical mass tracks of 1, 2,

3, 5, 8, 20, and 40 M from the Padova models. We assume a solar metallicity, which is well-matched to the gas-phase abundance of HII regions in the disk of M83. We

ﬁnd that the optical counterparts have a range of masses from ∼3–40 M . Detected

donor stars which have masses >∼ 8M are classified as HMXBs. There are a few sources with masses between 3–7M , and these are classified as IMXBs. Sources with masses between 3–7M are visually inspected to see if they suffer from a significant amount of extinction (e.g., dust lanes or are located near the center of the galaxy).

If so, these sources could end up having masses >∼ 8M and are classiﬁed accordingly (e.g., X197, X210, X224, X246, etc. ).

85 −10 40 MO •

20 MO •

−5 8 MO •

5 MO •

V 3 MO • 2 MO •

0 1 MO •

−1 0 1 2 3 V−I

Figure 4-5: Potential donor stars in M83 compared to theoretical mass tracks from the Padova models at solar metallicity.

86 The classiﬁcation for each X–ray source is given in Table 4.1. Note that for each

X–ray source we only include information for the candidate donor star or cluster.

In the case that a XRB is located in a crowded region where we can not with any

conﬁdence identify the donor star or cluster, we include information on all stars and

clusters within 2σ of the X–ray source. Luckily, in all but a few of these cases, all potential donors for a given XRB have masses >∼ 8M , so they can be safely classiﬁed as HMXBs.

We also assign each X–ray source a ‘confidence’ flag, where 1 =high confidence,

2 =medium conﬁdence, and 3 =low conﬁdence (Table 4.1). X–ray sources with one

candidate donor within 2σ in a region of low extinction will be assigned a Flag=1.

LMXBs with a high background are given lower conﬁdences (Flag=2) since there is a

possibility of an intermediate mass star going undetected. Similarly, IMXBs detected

in regions with high backgrounds are given lower conﬁdences (Flag=2) since photom-

etry of high–mass stars in these regions could lead to higher magnitude estimates.

LMXBs are given a low conﬁdence (Flag=3) if a dim optical source is detected exactly

2σ from the X–ray source. HMXBs and IMXBs are given low conﬁdences (Flag=3)

if there are multiple optical sources within 2σ, we cannot identify the donor, and all

optical sources have vastly diﬀerent mass ranges. HMXBs and IMXBs are assigned a

Flag=2 if there are multiple optical sources within 2σ, we cannot identify the donor,

yet all optical sources fall within the same mass range. If the X–ray source is located

within the nucleus of the galaxy, we assign it a low conﬁdence (Flag=3).

4.3.2 Classiﬁcation Based on Parent Cluster Ages

Recent works have identiﬁed a number of X–ray binaries in star clusters (Rangelov

et al. 2012, Fall et al. 2005, Zezas et al. 2002; Kaaret et al. 2004). In these cases,

no donor star can be identiﬁed because of the severe crowding. However, the optical

photometry provides the ages of the clusters.

87 High–mass (>8 M ) and intermediate–mass (3–8 M ) stars have hydrogen burning lifetimes of ≈ 10 Myr and ≈ 300–400 Myr respectively. Therefore, XRBs found

in clusters >400 Myr must therefore be LMXBs since only low–mass stars at this age

will remain. N–body simulations show that clusters <10 Myr are too young to form

binary systems through dynamical interactions (Garofali et al. 2012). This, coupled

with the fact that the ﬁrst binaries to form primordially are high-mass, XRBs found in

clusters <10 Myr are more likely to be HMXBs. Clusters 10–400 Myr are suﬃciently

old enough that they no longer host the most massive of stars (>8 M ). At this age, the more massive of the remaining stars have had suﬃcient time to form binary pairs through dynamical interactions (Garofali et al. 2012). We therefore classify any

XRBs found in clusters 10–400 Myr as IMXBs.

We ﬁnd 16 (∼ 7%) X–ray sources in M83 which have an optical counterpart that is a star cluster rather than an individual star. Donor star cluster age estimates and the resulting classiﬁcation of its XRB are found in Table 4.1.

88 Table 4.1: M83 X–ray Point Source Properties

X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

43 13:36:47.27 -29:51:38.90 35.90± 35.41 LMXB 2 47 13:36:48.30 -29:51:37.00 35.99± 35.41 LMXB 2

48 13:36:48.74 -29:52:29.60 36.48± 35.56 21.899± 0.008 -0.015±0.011 0.436±0.013 HMXB 3 Donor M∼12 M 49 13:36:49.12 -29:52:58.40 37.74± 36.26 LMXB 1 50 13:36:49.19 -29:51:25.00 36.31± 35.46 LMXB 2

51 13:36:49.20 -29:53:3.00 37.56± 36.30 24.370± 0.025 1.512±0.062 2.746±0.027 HMXB 1 Donor M∼9 M 52 13:36:49.43 -29:50:14.80 36.42± 35.51 LMXB 2 54 13:36:49.92 -29:52:59.40 36.63± 35.62 23.408± 0.023 0.960±0.042 1.779±0.028 LMXB 1 Cluster log(τ)∼8.86 55 13:36:49.94 -29:55:12.70 36.83± 35.74 LMXB 2

89 56 13:36:50.00 -29:52:30.60 36.20± 35.52 22.793± 0.010 -0.243±0.013 0.462±0.017 HMXB 1 Donor M∼9 M

58 13:36:50.11 -29:53:20.10 36.34± 35.54 23.039± 0.016 0.190±0.025 1.495±0.019 HMXB 2 Donor M∼11 M

59 13:36:50.34 -29:52:33.30 36.11± 35.53 22.393± 0.010 -0.171±0.013 0.281±0.016 HMXB 2 Donor M∼10 M

62 13:36:50.66 -29:52:47.40 36.05± 35.51 23.947± 0.021 0.172±0.029 0.514±0.034 IMXB 1 Donor M∼7 M 64 13:36:50.91 -29:52:59.50 36.36± 35.57 22.968± 0.013 -0.136±0.018 0.301±0.021 LMXB 3 21.832± 0.007 0.938±0.013 1.588±0.009 Two Potential 66 13:36:51.10 -29:50:6.90 35.93± 35.34 HMXB 2 21.304± 0.005 0.043±0.008 0.471±0.009 Donors

68 13:36:51.21 -29:50:9.80 35.88± 35.34 22.321± 0.009 0.113±0.013 0.528±0.014 HMXB 3 Donor M∼10 M

71 13:36:51.47 -29:50:43.40 35.97± 35.38 22.540± 0.011 0.674±0.018 1.252±0.014 HMXB 3 Donor M∼10 M 21.990± 0.010 -0.011±0.014 0.655±0.016 Two Potential 72 13:36:51.52 -29:51:43.20 36.86± 35.69 HMXB 3 22.076± 0.010 -0.087±0.013 0.647±0.016 Donors

73 13:36:51.65 -29:53:34.90 37.64± 36.00 20.853± 0.005 0.086±0.007 1.214±0.006 HMXB 1 Donor M∼16 M

74 13:36:51.65 -29:50:24.90 36.74± 35.62 25.569± 0.050 1.661±0.235 3.023±0.054 IMXB 1 Donor M∼5 M 75 13:36:51.74 -29:54:29.00 36.03± 35.54 LMXB 1 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 76 13:36:51.75 -29:54:31.10 36.53± 35.72 LMXB 1 22.711± 0.011 0.142±0.016 0.441±0.018 Multiple 77 13:36:51.77 -29:52:28.40 36.02± 35.45 23.574± 0.017 0.386±0.029 1.060±0.024HMXB 3 Potential 23.174± 0.015 0.505±0.024 0.881±0.022 Donors

79 13:36:51.90 -29:53:9.40 36.18± 35.56 23.138± 0.014 -0.112±0.019 0.372±0.023 IMXB 2 Donor M∼8 M 80 13:36:52.12 -29:52:22.80 35.91± 35.40 LMXB 1 23.630± 0.017 0.069±0.024 0.573±0.026 Two Potential 82 13:36:52.17 -29:49:37.20 35.87± 35.34 IMXB 2 23.039± 0.013 0.052±0.018 0.581±0.021 Donors 84 13:36:52.28 -29:49:20.10 35.95± 35.34 LMXB 1 23.970± 0.020 0.096±0.028 0.564±0.034 Two Potential 85 13:36:52.33 -29:50:46.40 37.50± 35.97 IMXB 2 24.136± 0.023 0.085±0.033 0.734±0.037 Donors 90 86 13:36:52.35 -29:53:1.20 36.48± 35.63 23.976± 0.019 1.897±0.053 2.648±0.021 IMXB 3 Donor M∼7 M

87 13:36:52.39 -29:52:51.60 36.61± 35.68 23.879± 0.018 -0.060±0.025 0.209±0.030 IMXB 3 Donor M∼7 M 88 13:36:52.41 -29:51:41.90 36.26± 35.49 LMXB 1

90 13:36:52.58 -29:55:31.60 36.35± 35.61 23.685± 0.016 1.724±0.042 2.056±0.019 HMXB 1 Donor M∼9 M

91 13:36:52.61 -29:51:46.90 36.32± 35.52 20.339± 0.004 0.275±0.006 0.757±0.006 HMXB 2 Donor M∼20 M 25.630± 0.038 -0.577±0.049 **** Two Potential 92 13:36:52.67 -29:49:37.40 36.18± 35.40 IMXB 2 24.783± 0.034 -0.208±0.045 0.245±0.058 Donors

94 13:36:52.81 -29:53:27.70 36.17± 35.53 23.623± 0.020 -0.148±0.028 0.591±0.032 IMXB 2 Donor M∼7 M 21.517± 0.008 -0.124±0.011 0.461±0.013 Two Potential 96 13:36:52.82 -29:51:43.60 36.13± 35.46 HMXB 2 21.304± 0.006 0.121±0.008 0.504±0.009 Donors 98 13:36:52.88 -29:51:37.60 36.51± 35.56 20.479± 0.005 0.671±0.009 1.131±0.006 LMXB 1 Cluster log(τ)∼8.96

99 13:36:52.88 -29:53:16.20 37.21± 35.86 22.948± 0.014 -0.138±0.019 0.266±0.023 HMXB 2 Donor M∼9 M 101 13:36:53.02 -29:53:10.10 36.33± 35.59 LMXB 1

102 13:36:53.13 -29:50:2.50 36.30± 35.46 24.842± 0.030 -0.050±0.041 -0.978±0.055 IMXB 3 Donor M∼6 M X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 103 13:36:53.14 -29:54:30.20 36.90± 35.72 21.416± 0.008 1.481±0.018 2.219±0.009 Quasar Candidate 1 109 13:36:53.51 -29:48:21.60 36.26± 35.49 LMXB 1

111 13:36:53.61 -29:53:36.50 37.73± 36.09 24.921± 0.027 5.051±0.069 2.922±0.030 IMXB 1 Donor M∼6 M 112 13:36:53.66 -29:51:54.50 36.05± 35.34 LMXB 1

113 13:36:53.68 -29:54:47.40 35.99± 35.45 24.363± 0.023 0.872±0.043 2.005±0.027 IMXB 1 Donor M∼7 M 21.800± 0.009 -0.125±0.012 0.714±0.015 22.567± 0.011 0.589±0.019 1.268±0.015 22.832± 0.013 -0.070±0.018 0.720±0.022 Multiple 114 13:36:53.81 -29:53:8.30 36.40± 35.59 21.740± 0.010 0.116±0.015 2.205±0.011HMXB 3 Potential 24.816± 0.024 -0.703±0.031 **** Donors 23.043± 0.016 0.600±0.033 1.212±0.023

91 23.970± 0.022 -0.276±0.028 -0.021±0.040

115 13:36:53.82 -29:51:1.10 36.00± 35.34 22.643± 0.012 0.076±0.017 0.979±0.018 IMXB 3 Donor M∼3 M

117 13:36:53.92 -29:51:14.60 37.37± 35.91 23.993± 0.028 0.908±0.061 2.944±0.029 IMXB 2 Donor M∼7 M

118 13:36:54.09 -29:49:33.10 36.43± 35.51 22.452± 0.010 -0.237±0.013 0.206±0.018 HMXB 1 Donor M∼10 M 23.741± 0.017 -0.269±0.023 0.053±0.029 23.175± 0.017 -0.157±0.023 0.745±0.025 Multiple 120 13:36:54.21 -29:53:7.30 36.39± 35.57 24.381± 0.023 -0.212±0.031 0.516±0.039HMXB 3 Potential 23.898± 0.019 -0.258±0.025 0.981±0.028 Donors 22.303± 0.010 -0.097±0.013 0.213±0.016 122 13:36:54.31 -29:51:43.70 36.08± 35.32 LMXB 2 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

22.837± 0.012 0.316±0.019 1.052±0.016 Multiple 23.302± 0.016 0.652±0.029 1.318±0.021 Potential 123 13:36:54.38 -29:52:57.30 36.42± 35.52 IMXB 3 22.942± 0.019 0.162±0.028 0.969±0.030 Donors 22.792± 0.013 0.139±0.019 0.347±0.021

124 13:36:54.55 -29:50:26.70 36.06± 35.38 23.682± 0.022 0.781±0.045 2.093±0.025 HMXB 2 Donor M∼9 M 125 13:36:54.66 -29:55:19.70 36.21± 35.54 LMXB 1

126 13:36:54.79 -29:53:9.00 36.34± 35.54 20.954± 0.007 -0.139±0.009 0.416±0.011 HMXB 2 Donor M∼15 M 130 13:36:55.05 -29:52:48.50 35.68± 35.28 LMXB 2 132 13:36:55.08 -29:53:15.40 35.97± 35.43 19.141± 0.003 0.236±0.005 1.190±0.004 IMXB 3 Cluster log(τ)∼7.63 133 13:36:55.09 -29:53:10.90 36.15± 35.46 LMXB 2 135 13:36:55.20 -29:54:3.10 36.74± 35.65 23.123± 0.020 0.158±0.028 0.563±0.030 IMXB 1 Cluster log(τ)∼8.3

92 137 13:36:55.37 -29:48:40.70 36.04± 35.43 LMXB 1 138 13:36:55.48 -29:55:9.80 37.83± 36.14 LMXB 1

139 13:36:55.56 -29:53:3.50 36.99± 35.76 24.075± 0.021 1.602±0.054 3.506±0.022 HMXB 1 Donor M∼8 M

143 13:36:56.38 -29:48:17.70 36.78± 35.75 22.643± 0.011 **** 2.060±0.013 HMXB 1 Donor M∼14 M

144 13:36:56.45 -29:52:57.60 35.88± 35.34 23.344± 0.019 0.291±0.028 1.423±0.028 IMXB 2 Donor M∼7 M

145 13:36:56.64 -29:49:12.40 38.05± 36.25 24.373± 0.023 -0.127±0.032 0.865±0.036 IMXB 1 Donor M∼6 M 146 13:36:56.66 -29:48:19.20 36.29± 35.48 LMXB 1 25.236± 0.034 -0.231±0.050 0.312±0.067 Two Potential 147 13:36:56.69 -29:53:21.30 36.39± 35.51 IMXB 2 24.782± 0.028 1.659±0.206 1.899±0.034 Donors 24.370± 0.021 1.312±0.042 0.876±0.028 Two Potential 148 13:36:56.79 -29:53:16.40 36.26± 35.45 IMXB 2 23.958± 0.021 0.729±0.037 1.623±0.027 Donors 150 13:36:57.25 -29:51:47.20 36.33± 35.43 LMXB 2

152 13:36:57.27 -29:53:39.40 38.08± 36.26 23.754± 0.020 1.606±0.051 2.448±0.022 HMXB 1 Donor M∼10 M

153 13:36:57.29 -29:50:31.90 36.33± 35.43 23.760± 0.076 0.274±0.079 0.758±0.080 IMXB 1 Donor M∼7 M X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

154 13:36:57.37 -29:50:35.40 36.47± 35.49 23.542± 0.070 0.334±0.074 0.938±0.073 IMXB 3 Donor M∼7 M 157 13:36:57.69 -29:50:39.20 36.38± 35.45 LMXB 1

158 13:36:57.80 -29:50:42.70 36.63± 35.54 23.368± 0.055 1.721±0.065 1.843±0.056 HMXB 1 Donor M∼11 M

160 13:36:57.91 -29:49:23.50 36.97± 35.72 24.498± 0.027 0.699±0.044 2.182±0.032 IMXB 1 Donor M∼7 M 23.134± 0.056 0.140±0.059 0.570±0.059 Two Potential 161 13:36:58.19 -29:51:24.10 36.80± 35.63 IMXB 2 23.772± 0.078 -0.030±0.081 0.384±0.084 Donors

162 13:36:58.23 -29:52:15.80 36.62± 35.57 23.673± 0.018 -0.372±0.023 -2.142±0.033 HMXB 1 Donor M∼37 M

163 13:36:58.29 -29:48:33.30 37.68± 36.07 22.254± 0.009 0.164±0.013 0.603±0.014 HMXB 2 Donor M∼9 M 24.239± 0.076 -0.170±0.079 -0.632±0.081 22.300± 0.035 0.222±0.036 0.693±0.037 Multiple 164 13:36:58.34 -29:51:27.50 35.88± 35.26 23.031± 0.051 -0.041±0.053 -0.033±0.056HMXB 3 Potential

93 24.185± 0.075 -0.174±0.078 0.253±0.081 Donors 22.250± 0.032 0.009±0.033 1.991±0.033 165 13:36:58.38 -29:51:4.70 37.55± 36.00 AGN 1

167 13:36:58.61 -29:53:38.90 36.08± 35.41 22.391± 0.010 -0.087±0.013 0.633±0.016 HMXB 1 Donor M∼11 M 168 13:36:58.63 -29:52:46.40 37.91± 36.17 LMXB 2 170 13:36:58.66 -29:51:6.70 36.28± 35.41 LMXB 2

173 13:36:58.80 -29:48:31.70 36.09± 35.41 24.619± 0.031 0.898±0.059 2.490±0.034 IMXB 1 Donor M∼6 M

175 13:36:58.91 -29:50:38.70 35.76± 35.23 24.106± 0.075 0.021±0.078 0.430±0.081 IMXB 1 Donor M∼7 M 176 13:36:58.91 -29:52:25.60 35.81± 35.30 LMXB 2

177 13:36:58.96 -29:50:25.00 36.00± 35.28 23.838± 0.073 0.480±0.077 0.451±0.079 IMXB 1 Donor M∼7 M 24.245± 0.021 0.682±0.049 1.812±0.025 Multiple 22.323± 0.010 0.389±0.016 0.785±0.014 Potential 178 13:36:59.09 -29:53:36.20 36.73± 35.62 HMXB 3 24.469± 0.023 3.339±0.079 2.394±0.026 Donors 23.836± 0.020 -0.316±0.026 0.644±0.032 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

179 13:36:59.11 -29:54:27.70 36.05± 35.45 22.947± 0.013 0.485±0.021 0.912±0.019 HMXB 3 Donor M∼9 M

180 13:36:59.19 -29:51:44.80 36.03± 35.45 23.795± 0.059 -0.139±0.061 1023.695±0.064 HMXB 1 Donor M∼11 M 182 13:36:59.32 -29:53:17.50 35.88± 35.32 LMXB 1 185 13:36:59.45 -29:49:59.00 38.57± 36.50 LMXB 2 187 13:36:59.58 -29:54:13.70 37.42± 35.93 LMXB 2 23.560± 0.058 0.344±0.061 **** Mutiple 22.810± 0.055 0.164±0.058 -0.101±0.059 Potential 188 13:36:59.65 -29:51:55.80 36.37± 35.65 IMXB 3 23.144± 0.054 0.164±0.056 0.712±0.058 Donors 24.344± 0.061 0.253±0.066 1.909±0.063

189 13:36:59.65 -29:52:22.70 36.09± 35.36 21.990± 0.028 0.206±0.029 0.640±0.030 HMXB 1 Donor M∼11 M

190 13:36:59.67 -29:51:8.80 36.48± 35.49 24.015± 0.072 -0.073±0.075 0.019±0.078 IMXB 1 Donor M∼7 M

94 24.216± 0.102 1.102±0.116 2.118±0.103 Two Potential 191 13:36:59.74 -29:52:16.10 35.91± 35.40 HMXB 3 24.163± 0.077 0.655±0.082 1.309±0.079 Donors 192 13:36:59.75 -29:51:54.80 36.44± 35.67 18.505±0.008 0.0719±0.008 0.815±0.008 IMXB 3 Cluster log(τ)∼7.38

193 13:36:59.77 -29:52:5.40 38.07± 36.27 23.529± 0.066 0.593±0.070 1.778±0.068 IMXB 3 Donor M∼7 M 23.007± 0.054 0.329±0.057 1.256±0.056 22.910± 0.055 1.361±0.064 2.620±0.055 Multiple 194 13:36:59.82 -29:52:2.40 36.85± 35.79 22.653± 0.040 0.404±0.042 0.881±0.043HMXB 3 Potential 23.974± 0.067 0.296±0.071 0.820±0.072 Donors 22.619± 0.046 0.371±0.048 1.036±0.048

196 13:36:59.83 -29:52:20.30 36.32± 35.48 23.566± 0.054 0.555±0.058 0.381±0.057 IMXB 1 Donor M∼7 M X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

197 13:36:59.98 -29:51:57.40 36.76± 35.83 22.061± 0.038 0.244±0.040 1.703±0.039 HMXB 2 Donor M∼13 M 22.455± 0.036 0.094±0.037 -0.113±0.038 22.609± 0.062 0.604±0.067 2.160±0.063 24.588± 0.057 1.341±0.060 2.488±0.059 Multiple 198 13:36:59.99 -29:51:50.30 37.87± 36.16 22.222± 0.034 0.277±0.036 0.072±0.037HMXB 3 Potential 21.951± 0.034 -0.030±0.035 0.383±0.036 Donors 23.689± 0.059 -0.583±0.061 **** 22.329± 0.036 0.058±0.038 0.961±0.038

200 13:37:0.04 -29:52:19.60 36.51± 35.57 21.341± 0.021 0.628±0.022 0.958±0.022 HMXB 2 Donor M∼16 M 201 13:37:0.06 -29:51:37.90 36.71± 35.63 LMXB 1

203 13:37:0.09 -29:53:29.60 37.13± 35.79 24.299± 0.034 0.325±0.049 1.546±0.043 IMXB 1 Donor M∼7 M 95 204 13:37:0.10 -29:51:45.70 37.14± 35.83 22.630± 0.057 0.325±0.061 1.261±0.058 HMXB 3 Donor M∼9 M 206 13:37:0.22 -29:51:52.40 37.38± 35.98 19.052± 0.008 1.413 ±0.009 1.968 ±0.008 LMXB 3 Cluster log(τ)∼10.02

208 13:37:0.24 -29:52:11.60 36.78± 35.71 22.151± 0.030 0.093±0.031 0.566±0.032 HMXB 2 Donor M∼11 M

209 13:37:0.28 -29:51:49.90 37.21± 35.92 21.585± 0.026 0.761±0.028 0.944±0.027 HMXB 3 Donor M∼13 M 22.424± 0.040 1.934±0.048 2.475±0.040 Multiple 22.774± 0.040 1.144±0.046 1.471±0.041 Potential 210 13:37:0.28 -29:51:51.30 36.91± 35.86 HMXB 2 22.962± 0.060 2.033±0.074 2.410±0.061 Donors 21.867± 0.031 0.892±0.034 1.249±0.032

213 13:37:0.35 -29:51:52.70 37.44± 36.01 21.475± 0.024 1.215±0.027 1.868±0.025 HMXB 2 Donor M∼20 M

214 13:37:0.37 -29:51:44.40 36.63± 35.60 22.986± 0.050 1.008±0.055 1.516±0.051 HMXB 3 Donor M∼10 M 215 13:37:0.37 -29:53:23.10 36.34± 35.45 LMXB 1 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

17.099± 0.006 0.006±0.006 0.818±0.006 Multiple 18.310± 0.007 0.086±0.007 0.525±0.007 Potential 216 13:37:0.42 -29:51:59.80 37.99± 36.24 16.813± 0.004 0.092±0.004 0.901±0.004HMXB 2 Donors 17.560± 0.006 -0.061±0.006 0.265±0.006 16.901±0.004 -0.002±0.004 0.437±0.004 Cluster log(τ)∼6.56

217 13:37:0.43 -29:50:54.10 36.77± 35.61 24.640± 0.084 2.662±0.107 2.481±0.085 IMXB 2 Donor M∼7 M 22.485± 0.048 0.048±0.050 0.857±0.050 Two Potential 218 13:37:0.44 -29:52:10.60 36.98± 35.79 HMXB 3 22.368± 0.052 0.131±0.054 0.852±0.054 Donors 219 13:37:0.44 -29:52:23.00 35.95± 35.32 LMXB 1

220 13:37:0.48 -29:51:55.90 37.93± 36.20 19.174± 0.008 0.166±0.008 0.904±0.008 HMXB 2 Donor M∼33 M 222 13:37:0.55 -29:52:10.30 36.96± 35.78 19.436± 0.010 0.258±0.010 0.725±0.010 IMXB 3 Cluster log(τ)∼7.68 96 223 13:37:0.55 -29:51:59.20 37.85± 36.18 18.943± 0.008 0.468±0.009 0.911±0.008 HMXB 3 Donor M∼38 M

224 13:37:0.57 -29:52:4.00 37.53± 36.04 20.556± 0.016 0.053±0.016 0.095±0.017 HMXB 3 Donor M∼21 M 225 13:37:0.60 -29:51:46.90 37.33± 35.90 LMXB 1

226 13:37:0.64 -29:51:53.60 37.11± 35.86 22.205± 0.028 -0.372±0.029 0.781±0.030 HMXB 2 Donor M∼11 M 23.476± 0.018 0.284±0.027 0.634±0.026 Multiple 23.968± 0.019 0.269±0.028 1.041±0.027 Potential 227 13:37:0.64 -29:53:19.70 38.12± 36.27 IMXB 2 24.456± 0.027 0.478±0.042 1.109±0.038 Donors 24.267± 0.023 1.308±0.053 2.066±0.027

228 13:37:0.68 -29:52:6.30 37.81± 36.14 21.284± 0.032 0.385±0.034 0.896±0.033 HMXB 2 Donor M∼14 M X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

20.676± 0.016 -0.175±0.016 0.530±0.017 20.273± 0.018 -0.005±0.019 0.473±0.019 19.535± 0.011 0.151±0.011 0.555±0.011 18.820± 0.010 0.064±0.010 0.750±0.010 Cluster log(τ)∼7.38 229 13:37:0.68 -29:52:0.50 37.63± 36.09 19.909± 0.012 0.421±0.013 0.863±0.013HMXB 2 Multiple 19.688± 0.012 -0.148±0.012 0.767±0.013 Potential 21.285± 0.021 0.072±0.022 **** Donors 19.370± 0.012 0.163±0.013 0.786±0.013 20.366± 0.020 0.366±0.022 1.134±0.021

230 13:37:0.69 -29:52:4.30 37.47± 36.03 20.505± 0.015 0.136±0.016 0.807±0.016 HMXB 2 Donor M∼20 M

231 13:37:0.71 -29:51:56.80 37.78± 36.14 18.608± 0.008 0.029±0.008 0.285±0.009 HMXB 2 Donor M∼43 M

97 20.885± 0.019 0.052±0.020 0.860±0.020 Two Potential 232 13:37:0.89 -29:51:58.30 37.63± 36.08 HMXB 3 23.667± 0.049 0.107±0.052 -0.710±0.050 Donors 21.228± 0.021 0.409±0.022 1.433±0.022 19.513± 0.014 0.318±0.015 0.899±0.015 Multiple 19.392± 0.012 0.017±0.012 0.970±0.013 Potential 234 13:37:0.95 -29:52:2.80 38.34± 36.39 HMXB 3 21.604± 0.018 -1.024±0.019 **** Donors 18.801± 0.007 0.577±0.008 0.904±0.007 19.443± 0.013 0.136±0.014 -0.945±0.016 236 13:37:1.06 -29:52:45.50 37.75± 36.10 LMXB 2 237 13:37:1.11 -29:51:52.10 37.58± 36.03 LMXB 3

239 13:37:1.14 -29:51:55.20 37.40± 35.97 21.473± 0.034 0.495±0.036 2.414±0.034 HMXB 2 Donor M∼16 M 23.051± 0.017 0.262±0.029 1.337±0.021 Multiple 240 13:37:1.17 -29:54:49.20 36.47± 35.54 23.648± 0.019 0.242±0.029 0.661±0.028HMXB 3 Potential 22.978± 0.014 0.173±0.019 1.191±0.021 Donors X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

241 13:37:1.19 -29:51:57.30 37.62± 36.05 19.781± 0.011 0.167±0.012 0.597±0.012 HMXB 3 Donor M∼26 M 242 13:37:1.26 -29:52:2.30 37.16± 35.85 LMXB 1 244 13:37:1.30 -29:51:36.50 37.38± 35.91 LMXB 2

245 13:37:1.31 -29:51:28.80 36.17± 35.40 20.478± 0.014 0.104±0.015 0.396±0.015 HMXB 3 Donor M∼19 M

246 13:37:1.46 -29:51:41.50 35.95± 35.36 23.619± 0.077 0.542±0.081 1.570±0.080 HMXB 2 Donor M∼9 M 247 13:37:1.46 -29:51:54.50 36.37± 35.54 21.316± 0.025 1.670±0.030 2.827±0.025 LMXB 3 Cluster log(τ)∼10.24 248 13:37:1.48 -29:53:26.50 38.39± 36.41 LMXB 2

251 13:37:1.63 -29:51:28.20 38.22± 36.32 21.617± 0.023 -0.032±0.024 0.155±0.025 HMXB 3 Donor M∼13 M

252 13:37:1.63 -29:47:42.80 37.92± 36.21 24.833± 0.032 1.176±0.064 1.967±0.038 IMXB 1 Donor M∼6 M

254 13:37:1.68 -29:52:11.30 35.95± 35.36 23.434± 0.062 0.182±0.065 0.723±0.065 IMXB 2 Donor M∼7 M 257 13:37:1.84 -29:51:19.10 35.88± 35.30 20.102± 0.016 0.201±0.017 1.047±0.016 IMXB 2 Cluster log(τ)∼7.53 98 258 13:37:2.01 -29:55:17.90 37.98± 36.22 25.581± 0.041 1.711±0.090 2.509±0.047 IMXB 1 Donor M∼4 M 259 13:37:2.15 -29:55:6.00 36.98± 35.75 LMXB 1 260 13:37:2.20 -29:51:44.10 36.31± 35.45 LMXB 2

264 13:37:2.34 -29:52:6.30 35.93± 35.32 23.789± 0.060 0.282±0.063 1.147±0.063 IMXB 1 Donor M∼7 M 22.990± 0.062 0.510±0.065 0.983±0.065 21.379± 0.021 0.157±0.022 0.568±0.022 Multiple 265 13:37:2.42 -29:51:26.10 37.24± 35.83 20.932± 0.018 1.070±0.020 1.261±0.018HMXB 2 Potential 23.074± 0.072 0.107±0.075 0.057±0.078 Donors 20.017± 0.012 0.433±0.013 1.240±0.012

266 13:37:2.44 -29:51:19.60 35.76± 35.23 20.906± 0.018 0.078±0.019 0.370±0.019 HMXB 1 Donor M∼14 M

267 13:37:2.45 -29:53:19.20 37.41± 35.92 24.091± 0.025 1.327±0.052 2.470±0.028 IMXB 1 Donor M∼7 M 269 13:37:2.67 -29:48:24.40 36.33± 35.48 21.626± 0.011 0.152±0.016 0.951±0.016 IMXB 3 Cluster log(τ)∼8.51 270 13:37:2.83 -29:51:40.50 35.95± 35.36 LMXB 2 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

23.182± 0.018 -0.061±0.024 0.596±0.028 Two Potential 273 13:37:3.11 -29:55:31.50 36.00± 35.51 IMXB 2 23.468± 0.019 0.150±0.027 0.586±0.029 Donors

274 13:37:3.29 -29:52:26.70 37.16± 35.81 22.331± 0.034 1.420±0.039 1.639±0.035 HMXB 1 Donor M∼13 M 276 13:37:3.48 -29:51:42.90 35.73± 35.20 LMXB 2 277 13:37:3.50 -29:53:31.20 36.33± 35.43 LMXB 1

278 13:37:3.56 -29:53:19.90 36.05± 35.32 23.190± 0.014 0.143±0.021 0.793±0.021 IMXB 1 Donor M∼8 M 280 13:37:3.84 -29:52:22.90 35.88± 35.26 LMXB 1

281 13:37:3.88 -29:49:30.40 37.97± 36.20 23.666± 0.017 -0.021±0.024 0.183±0.030 IMXB 2 Donor M∼7 M 282 13:37:3.93 -29:53:22.00 35.68± 35.23 LMXB 1 24.031± 0.022 0.194±0.031 0.608±0.033 Two Potential 283 13:37:4.18 -29:53:12.20 36.20± 35.38 IMXB 2 23.841± 0.021 0.334±0.031 1.040±0.032 Donors

99 22.960± 0.014 0.286±0.021 1.151±0.021 Multiple 23.823± 0.021 0.457±0.032 1.087±0.028 Potential 284 13:37:4.28 -29:54:3.60 38.49± 36.46 HMXB 3 22.891± 0.012 0.655±0.020 1.257±0.016 Donors 22.533± 0.011 1.841±0.029 2.116±0.013

285 13:37:4.37 -29:51:30.70 37.34± 35.88 23.129± 0.047 -0.062±0.049 0.327±0.050 HMXB 1 Donor M∼8 M

286 13:37:4.37 -29:51:21.50 38.66± 36.54 23.836± 0.066 0.108±0.069 0.662±0.070 HMXB 3 Donor M∼8 M

289 13:37:4.56 -29:51:7.80 35.90± 35.26 23.383± 0.055 0.052±0.058 -0.096±0.060 HMXB 1 Donor M∼9 M 290 13:37:4.65 -29:51:20.40 36.77± 35.61 20.283± 0.018 0.761±0.019 1.399±0.018 LMXB 1 Cluster log(τ)∼8.91 291 13:37:4.65 -29:50:54.50 36.36± 35.46 LMXB 2

293 13:37:4.70 -29:50:40.20 35.97± 35.30 23.694± 0.069 -0.072±0.072 0.099±0.076 HMXB 1 Donor M∼8 M 294 13:37:4.73 -29:50:58.90 35.94± 35.32 19.836± 0.017 0.100±0.018 0.463±0.018 HMXB 2 Cluster log(τ)∼6.48 295 13:37:4.74 -29:48:52.00 37.18± 35.82 LMXB 1

296 13:37:4.83 -29:51:7.40 35.85± 35.26 23.448± 0.057 0.414±0.061 1.273±0.059 HMXB 3 Donor M∼8 M X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

21.984± 0.008 0.091±0.011 0.761±0.013 Two Potential 298 13:37:4.90 -29:53:39.50 36.02± 35.34 HMXB 3 26.768± 0.047 **** 5.932±0.048 Donors 299 13:37:5.13 -29:52:7.10 39.35± 36.89 LMXB 2 300 13:37:5.15 -29:53:17.10 35.74± 35.26 LMXB 1

301 13:37:5.26 -29:51:21.80 36.27± 35.38 24.269± 0.085 0.941±0.096 1.583±0.086 HMXB 3 Donor M∼9 M 302 13:37:5.41 -29:54:9.60 35.87± 35.34 23.919± 0.022 1.462±0.048 2.280±0.025 LMXB 3

303 13:37:5.47 -29:52:34.00 37.37± 35.91 24.769± 0.098 -0.182±0.102 0.563±0.104 IMXB 1 Donor M∼6 M 304 13:37:5.50 -29:50:32.60 36.01± 35.32 LMXB 1 307 13:37:5.76 -29:49:23.20 35.69± 35.23 LMXB 1 308 13:37:5.84 -29:48:21.90 36.07± 35.43 LMXB 1

309 13:37:5.99 -29:51:59.10 36.63± 35.54 24.257± 0.075 1.396±0.087 2.399±0.076 HMXB 1 Donor M∼9 M 100 312 13:37:6.19 -29:52:32.00 36.43± 35.46 27.776± 0.170 **** 4.240±0.171 IMXB 2 Donor M∼4 M

314 13:37:6.61 -29:49:44.20 36.01± 35.28 21.192± 0.006 -0.070±0.008 0.341±0.009 HMXB 1 Donor M∼16 M 315 13:37:6.63 -29:51:7.60 35.84± 35.23 LMXB 1

317 13:37:6.78 -29:50:57.80 36.41± 35.48 22.491± 0.035 0.108±0.036 0.437±0.037 HMXB 1 Donor M∼10 M

318 13:37:6.97 -29:49:34.60 35.99± 35.32 21.851± 0.007 0.268±0.011 0.637±0.011 HMXB 3 Donor M∼12 M

321 13:37:7.11 -29:51:1.60 38.43± 36.43 22.771± 0.053 0.050±0.055 -0.047±0.057 HMXB 2 Donor M∼10 M 322 13:37:7.12 -29:52:2.10 36.59± 35.52 21.713± 0.042 0.800±0.046 1.490±0.043 LMXB 1 Cluster log(τ)∼8.76 323 13:37:7.15 -29:51:8.80 35.76± 35.26 LMXB 1 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

22.563± 0.013 0.367±0.021 0.798±0.018 22.695± 0.015 0.241±0.023 0.979±0.021 21.642± 0.007 0.084±0.011 0.483±0.011 Multiple 327 13:37:7.52 -29:48:59.40 36.72± 35.61 22.013± 0.009 -0.077±0.013 0.213±0.015HMXB 3 Potential 22.537± 0.015 -0.091±0.021 0.245±0.026 Donors 23.262± 0.014 -0.013±0.021 0.209±0.024 23.623± 0.017 -0.028±0.025 0.380±0.029

328 13:37:7.56 -29:49:18.80 36.35± 35.46 23.980± 0.021 0.359±0.034 1.013±0.029 IMXB 3 Donor M∼7 M 23.331± 0.057 -0.089±0.059 0.288±0.061 Multiple 329 13:37:7.68 -29:50:56.90 35.89± 35.30 22.698± 0.036 -0.136±0.037 0.386±0.038HMXB 2 Potential 23.042± 0.046 -0.196±0.048 0.067±0.050 Donors 101 331 13:37:7.94 -29:52:14.30 35.72± 35.20 24.386± 0.120 0.454±0.127 2.995±0.120 IMXB 1 Donor M∼7 M 332 13:37:8.19 -29:49:16.80 36.52± 35.53 22.499± 0.014 0.112±0.021 0.685±0.021 IMXB 2 Cluster log(τ)∼8.06 24.733± 0.106 2.536±0.120 2.256±0.107 Multiple 333 13:37:8.21 -29:52:57.40 35.83± 35.28 24.018± 0.077 0.127±0.080 0.744±0.085IMXB 2 Potential 24.314± 0.128 0.828±0.139 2.710±0.129 Donors 335 13:37:8.25 -29:51:25.00 35.92± 35.30 LMXB 2 23.187± 0.046 -0.100±0.048 0.018±0.050 Multiple 20.929± 0.020 -0.071±0.021 0.048±0.022 Potential 337 13:37:8.36 -29:52:55.60 35.92± 35.30 HMXB 2 20.887± 0.022 0.060±0.023 0.717±0.023 Donors 20.890± 0.026 0.068±0.027 0.751±0.027 343 13:37:8.79 -29:52:6.30 35.67± 35.20 LMXB 1 344 13:37:8.79 -29:53:34.90 36.24± 35.48 LMXB 1

345 13:37:9.05 -29:49:38.50 36.78± 35.63 24.274± 0.022 0.274±0.035 1.005±0.032 IMXB 1 Donor M∼6 M

346 13:37:9.07 -29:51:25.10 35.86± 35.28 25.264± 0.134 2.551±0.200 3.217±0.135 IMXB 3 Donor M∼4 M X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag)

347 13:37:9.20 -29:49:24.80 35.86± 35.34 21.943± 0.008 1.666±0.020 2.762±0.009 HMXB 2 Donor M∼18 M 349 13:37:9.87 -29:50:26.80 36.18± 35.38 LMXB 1

351 13:37:10.31 -29:49:18.40 36.10± 35.40 23.863± 0.019 -0.242±0.025 0.047±0.033 HMXB 2 Donor M∼10 M 23.916± 0.021 1.105±0.045 2.253±0.024 Two Potential 358 13:37:11.88 -29:52:15.50 36.95± 35.71 IMXB 2 23.473± 0.018 0.205±0.025 1.116±0.026 Donors

363 13:37:12.45 -29:51:40.10 36.93± 35.69 23.789± 0.019 0.729±0.031 1.339±0.024 IMXB 1 Donor M∼7 M 24.324± 0.027 0.653±0.052 1.370±0.034 Two Potential 366 13:37:12.53 -29:51:54.90 37.59± 36.01 IMXB 2 24.354± 0.030 -0.065±0.041 0.789±0.047 Donors 24.654± 0.034 0.265±0.049 0.616±0.063 Two Potential 370 13:37:13.13 -29:52:38.20 36.74± 35.61 IMXB 2 25.003± 0.035 -0.074±0.046 0.448±0.059 Donors

371 13:37:13.65 -29:52:0.10 35.84± 35.26 23.122± 0.014 -0.060±0.020 0.546±0.024 HMXB 1 Donor M∼10 M 102 24.739± 0.032 -0.003±0.045 1.449±0.047 Multiple 24.313± 0.026 1.771±0.070 2.439±0.029 Potential 377 13:37:14.40 -29:51:30.30 36.04± 35.34 IMXB 2 24.965± 0.032 1.379±0.109 2.730±0.035 Donors 26.562± 0.050 **** 3.502±0.053

378 13:37:14.45 -29:51:48.80 37.42± 35.93 21.040± 0.006 0.633±0.008 1.832±0.007 HMXB 1 Donor M∼18 M

381 13:37:15.01 -29:51:38.70 36.13± 35.36 24.794± 0.029 -0.040±0.040 1.787±0.041 IMXB 1 Donor M∼6 M

384 13:37:16.20 -29:52:2.20 36.85± 35.67 24.437± 0.028 0.180±0.038 0.575±0.044 IMXB 1 Donor M∼6 M 393 13:37:17.98 -29:52:11.80 36.02± 35.38 LMXB 1 395 13:37:18.42 -29:51:18.50 36.13± 35.41 LMXB 1 396 13:37:18.57 -29:52:7.60 35.90± 35.36 LMXB 1

397 13:37:18.89 -29:50:13.70 36.59± 35.62 23.697± 0.021 0.711±0.035 1.464±0.028 IMXB 3 Donor M∼7 M 402 13:37:19.64 -29:51:31.60 36.93± 35.75 21.594± 0.009 1.474±0.018 1.966±0.011 AGN 1 405 13:37:20.02 -29:51:30.50 35.88± 35.34 AGN 1 408 13:37:21.11 -29:52:42.60 36.16± 35.54 LMXB 1 X–ray R.A. Decl. log(LX) v (B − V )(V − I) Classiﬁcation Flag Notes ID (J2000) (J2000) (erg/s) (mag) (mag) (mag) 409 13:37:21.43 -29:51:21.30 36.43± 35.65 LMXB 1 412 13:37:22.15 -29:52:7.90 36.81± 35.79 22.714± 0.012 6.300± 0.012 3.250±0.013 AGN 1 103 4.3.3 Spatial Distribution of XRBs

Figure 4-6 shows the spatial locations of our classified X–ray sources. Sources we have classified as HMXBs, IMXBs, and LMXBs are identified by blue, green, and red circles respectively. All other classified X–ray sources such as background AGN and quasars are shown in magenta. Note that the central bulge region of M83 is dominated by HMXBs, and not LMXBs as is often assumed. We also find that the disk region is populated by both HMXBs and LMXBs. Similar to what we see in the previous chapter for M101, we find that LMXBs are found equally distributed throughout

M83. Ultimately, classifying XRBs based on their location within a galaxy proves to be extremely inaccurate.

4.4 X–ray Luminosity Function

Figure 4-7 compares the XLF from our individually classiﬁed HMXBs (blue),

IMXBs (green), and LMXBs (red). We do not find that the XLF for HMXBs extends to significantly higher fluxes (unlike in M101), although we do find that the XLF for HMXBs does have a flatter distribution than those for LMXBs and IMXBs. This

ﬂatter distribution for HMXBs relative to LMXBs is similar to what we found in

M101, and suggests that our new method for classifying the diﬀerent type of XRBs is working.

In this section, we present fits and quantitative results for the XLFs of our classified HMXBs, IMXBs, and LMXBs. We use the same approach used in Section 3.5 for M101. Figures 4-8 and 4-9 summarizes our results, and shows a simple power law fit to the binned distributions (left), a truncated power law fit to the cumulative luminosity distribution (middle), and the maximum–likelihood fitting of a Schechter function to the observed luminosity function (right).

104 Figure 4-6: The spatial distribution of X–ray sources in M83. Sources we have classified as HMXBs, IMXBs, and LMXBs are identified by blue, green, and red circles respectively. All other classified X–ray sources such as background AGN and quasars are shown in magenta.

105 Figure 4-7: Cumulative X–ray luminosity functions for classiﬁed HMXBs (blue), IMXBs (green), and LMXBs (red) in M83.

106 107

Figure 4-8: Statistical ﬁts to total XLF of XRBs in M83. The panels in Figure 4-8 show the results of the three methods when applied to the total X–ray luminosity function of all XRBs in M83 (HMXBs+IMXBs+LMXBs).

The left and center panels show that the overall XLF appears to be well ﬁt by a single power-law with α ≈ -1.4 to -1.5 which is somewhat ﬂatter than the α ≈-2 observed in M101. The middle panel shows no evidence of a Schechter-like downturn, since the statistic N0 =8.84±6.08 doesn’t indicate a detection at the ≥ 3σ level (only ∼1.5σ).

108 109

Figure 4-9: Statistical ﬁts to XLF of HMXBs (top), IMXBs (middle), and LMXBs (bottom) in M83. Figure 4-9 shows the ﬁtting results for HMXBs (top row), IMXBs (middle row),

and LMXBs (bottom row). Here we show ﬁts for sources at 1 + 2 conﬁdence levels

since we find little difference in the results when we only use higher confidence levels.

The X–ray luminosity functions for HMXBs are well ﬁt by a single power law with α=-

1.36±0.10 from the binned method and α=-1.47±0.05 from the cumulative one. There is no evidence for a downturn at the bright end. For IMXBs, the X–ray luminosity functions are slightly steeper and are ﬁt with a α=-1.37±0.11 (α=-1.51±0.04) from the binned (cumulative) method. There is no evidence of a Schechter-like downturn.

LMXBs are ﬁt with a α=-1.46±0.08 (α=-1.55±0.11) from the binned (cumulative) method and show no evidence for a downturn at the bright end, especially coupled with the right panel where the contours don’t close around any particular value, instead giving a lower X–ray ﬂux limit.

4.5 Discussion

Historically, the only locations where astronomers felt comfortable determining the XLF of LMXBs in late-type galaxies is in the bulges of spiral galaxies. The disks of late-type galaxies are the location for on-going star formation, and must therefore contain a mix of HMXBs and LMXBs. No work yet exists which considers or includes

LMXBs outside of these regions. To date, astronomers have only determined the XLF of LMXBs in late–type galaxies by using all X–ray sources within a given eﬀective radius of the bulge. However, we believe that many of these sources are likely HMXBs.

By using the results from our new method of diﬀerentiating between HMXBs and

LMXBs directly, we determined the XLF of LMXBs in star forming galaxies, across the entire galaxy for the ﬁrst time.

110 Like in M101, we dont see a Schechter-like downturn at the high luminosity end for any sub-population of XRBs. We need more statistics, i.e. more galaxies, to determine if one actually exists.

111 Chapter 5

Conclusions and Future Prospects

5.1 Conclusions

Relatively little work has previously been conducted to study XRB properties in late type galaxies. The work that has been done has often been clouded with large systematic uncertainties thought to be either inherent to the data or as a result of the XRB classiﬁcation method applied. For this dissertation, we present two new and more accurate methods for classifying XRBs from their optical properties, either inferred from their parent star cluster age or directly from their donor star. In this way, by classify XRBs individually rather than statistically based on their location within the galaxy or through the use of X–ray color–color diagrams, the hope is to eliminate the large uncertainties that exist in current studies of XRBs in star-forming galaxies.

Optical images taken with the Hubble Space Telescope allow us to reliably identify background galaxies/quasars, foreground stars, and supernova remnants, yield- ing low-contamination samples of HMXBs, IMXBs, and LMXBs. We used available or previously published catalogues of XRBs for the merging Antennae galaxies,

NGC 4449, M101, and M83, including positions and X–ray luminosities, based on data from the Chandra X–ray Observatory. We determined the level of astrometric

112 matching between the X–ray and optical positions, and used this to identify parent clusters to XRBs in all galaxies, and donor stars in the large spirals M101 and M83.

We were able to estimate the masses of the donor star in XRBs, either directly or from the age of the parent cluster. In this work, we deﬁned a LMXB as having a donor star mass <3 M , a IMXB as having a donor star mass 3–8M , and a HMXB as having a donor star mass >8 M . XRBs which are still embedded in their parent cluster don’t have resolved donor stars. In these cases, we instead use the hydrogen burning lifetimes predicted by stellar evolutionary models as a guide to classifying

XRBs. We conclude that parent clusters with ages <10 Myr host HMXBs, parent clusters with ages 10–400 Myr host IMXBs, and parent clusters with ages >400 Myr host LMXBs.

In Chapter 2 we studied the properties of clusters which host XRBs in the merging Antennae galaxies and in the starburst NGC 4449. We found that the masses of clusters hosting HMXBs are consistent with being drawn from the overall mass and density distributions of the clusters. We also find tentative evidence that intermediate (log(τ/yr)=8–8.6) parent clusters which likely host IMXBs are weighted towards higher masses and densities. If confirmed, these differences imply that HMXBs and

IMXBs have diﬀerent formation mechanics, with dynamical interactions playing a far more important role in forming IMXBs and LMXBs than HMXBs.

The luminosity function is a basic tool used to characterize populations of XRBs in nearby galaxies. The shape of the XLFs of both LMXBs and HMXBs are believed to be well described by a Schechter-like distribution, with a power law component, and a downturn or exponential cutoﬀ at the high luminosity end. It has been suggested that the power-law indices and break luminosities diﬀer between LMXBs and

HMXBs. Of particular interest are breaks around the Eddington luminosity limit

∼2–5×1038 erg s−1 for an accreting neutron star (Sarazin, Irwin & Bregman 2000).

113 We have used our method of separating HMXBs from LMXBs based on the mass of the donor star in the spirals M101 and M83, and found that the XLFs for these different populations are different, with a shallower power-law index for the HMXBs than for LMXBS, as has been found previously in the literature. This difference suggests that our new method works. Further, the differences indicate that HMXBs and LMXBs have different formation mechanisms. We believe that this is the first time that XLFs have been presented for a population of IMXBs. The IMXBs appear to be quite similar to LMXBs, both when comparing their parent cluster properties and their XLFs. We do not see any statistically significant evidence for a cutoff at the high luminosity end at any population of XRBs in M101 or M83 .

Ultimately, the results from each of the projects described in this thesis paint a consistent picture. Speciﬁcally, that HMXBs are formed primordially and that IMXB and LMXBs are formed through dynamical capture. This is supported by the results of our statistical tests which show that IMXBs and not HMXBs are more likely to be found in more massive and more dense clusters. Additionally, HMXB XLFs are shallower in comparison to LMXB and IMXB XLFs, supporting the idea that these populations are formed through diﬀerent mechanisms.

5.2 Future work

Ultimately, better statistics will be required to characterize the shapes of X–ray luminosity functions for HMXBs, IMXBs, and LMXBs in star-forming galaxies, and to compare them with properties like the star formation rate and mass of the host.

In Table 5.1, we compile a sample of ∼ 50 star-forming galaxies (including M83 and

M101) located within ∼ 10 Mpc that have suﬃcient archival X–ray observations from

Chandra and optical observations from Hubble. The archival Hubble data is deep

114 enough that we can detect stars down to ∼ 3 M . The X–ray and optical data can be analyzed in the same way as presented in Chapters 2 through 4.

Using this sample of nearby star–forming galaxies in the future, we will produce catalogs of XRBs (cleaned of foreground stars and background AGN) in each galaxy.

Each XRB will be classiﬁed as a HMXB, IMXB, or LMXB, based on either the age of its parent cluster or the mass of its donor star. We will produce catalogs of ancient globular clusters in each galaxy, and estimate the extinction-corrected star formation rate (from archival far-ultraviolet images taken with GALEX combined with 24µm images taken with Spitzer), and total galaxy mass estimates from near-infrared Spitzer luminosities, all within the HST ﬁeld of view.

What is the scaling relation between LMXBs and galaxy mass in star-forming galaxies? A large amount of eﬀort has gone into determining the scaling relations between the XLF of LMXBs and galaxy mass (e.g., Gilfanov 2004). However, similar studies in star–forming galaxies have not been possible to date, because there hasn’t been a clean way to separate LMXBs from HMXBs using X–ray observations alone.

One of the main goals of future work will be to produce this relation for the ﬁrst time, and to compare with that found for early-type galaxies.

With the techniques outlined in this work applied to a large number of star– forming galaxies, we can perform statistical tests to establish whether a single or a double power law best ﬁts the XLF of LMXBs for each galaxy. Do these normalized distributions match the distribution of XRBs in early–type galaxies? If so, this indicates that the LMXB population in spiral galaxies are similar to the populations found in early–type galaxies. In addition, we will compare the shapes between diﬀerent galaxies to help determine how “universal” the XLF of LMXBs is.

115 Table 5.1: Galaxy Sample and Basic Properties

Galaxy Distance MB log(O/H)+12 Ref Name (Mpcs) (mag)

Starbursts

NGC 1569 2.2 -16.6 8.2 Leitherer et al. 2011

NGC 1705 5.1 -15.8 8.2 Leitherer et al. 2011

NGC 3125 11.5 -18.0 8.3 Chandar et al. 2004

NGC 4214 2.9 -17.2 8.1 Leitherer et al. 2011

NGC 4449 3.8 -18.2 8.3 Annibali et al. 2008

NGC 5253 3.6 -16.8 8.2 Leitherer et al. 2011

Spirals

M83 4.5 -19.8 8.7 Pilyugin et al. 2012

M51 7.6 -20.7 8.8 Pilyugin et al. 2004

M101 7.4 -21.1 8.9 Leitherer et al. 2011

NGC 628 7.3 -19.5 8.5 Pilyugin et al. 2004

NGC 3627 8.1 -21.2 9.2 Prieto et al. 2008

NGC 4736 4.9 -19.7 8.5 Pilyugin et al. 2004

NGC 5055 8.7 -20.5 9.0 Leitherer et al. 2011

NGC 7793 3.9 -18.5 8.9 Edmunds & Pagel 1984

NGC 4258 7.6 -20.1 8.6 Dack & McCall 2012

NGC 1433 8.3 -19.0 9.0 Florido et al. 2012

NGC 5068 6.6 -18.7 8.7 Bibby & Crowther 2012

NGC 6744 8.4 -20.2 8.5 Talent 1982

116 What role do globular star clusters play in producing LMXBs in late–type galax-

ies? In early–type galaxies, ≈4% of all globular star clusters host a LMXB, and the

fraction of LMXBs found in globular clusters ranges from ≈25-70%, depending on the galaxy. We will make separate composite globular cluster populations for the spirals and dwarf starbursts. For the spirals, based on our previous work in M51 and M101, we expect to detect between 30 and 70 globular clusters per galaxy.

How similar are the shapes of the luminosity function of HMXBs in diﬀerent star- forming galaxies? For a given galaxy, how do the XLFs of LMXBs and HMXBs compare? The shape of the XLF for HMXBs can generally be described by a power-

dN −α law, dL ∝ L , with a typical value of α ≈ -1.6, but with published values ranging from ≈ -1.2 to -2.5 for nearby star-forming galaxies. By producing clean samples of HMXBs for ≈ 50 galaxies, we will be able to deﬁnitively establish the shape and universality of this fundamental distribution.

Does the SFR set the normalization of the HMXB XLF? Current works suggest that the SFR is the main driver of the amplitude of the XLF, with metallicity possibly playing a secondary role. We will establish the most robust XLFs of HMXBs to date, with homogeneously determined SFRs, and assess the role of metallicity by doing so for both high and low metallicity galaxy samples.

To investigate the constraints on the formation of XRBs in star–forming galaxies, we need to establish where XRBs are created. Are the XRBs found in the ﬁeld and in globular clusters from the same population? If not, we expect ﬁeld XRBs to correlate with the host’s galaxy’s properties and globular cluster XRBs to correlate with globular cluster properties. We will compare the X–ray colors, luminosities, and

XLF of XRBs found in the ﬁeld to those found in globular clusters.

By applying the techniques that we have developed in this dissertation, a full characterization of the HMXB and LMXB populations in star-forming galaxies is now within reach.

117 A special thank you to Angus Mok for his invaluable help with our XLF ﬁtting. We also want to thank Blagoy Rangelov and Andrea Prestwich for their helpful discussions and insight.

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