# The Local Environment of Ultraluminous X-ray Sources Viewed by XMM Newton’s
Optical Monitor

C. T. Berghea and R. P. Dudik  United States Naval Observatory, Washington, DC
20392<EMAIL_ADDRESS><EMAIL_ADDRESS>J. Tincher
St. John’s College, Annapolis, MD 21401  L. M. Winter  Atmospheric and
Environmental Research, Lexington, MA 02421

###### Abstract

We have used XMM-Newton’s Optical Monitor (OM) images to study the local
environment of a sample of 27 Ultraluminous X-ray Sources (ULXs) in nearby
galaxies. UVW1 fluxes were extracted from 100 pc regions centered on the ULX
positions. We find that at least 4 ULXs (out of 10 published) have spectral
types that are consistent with previous literature values. In addition the
colors are similar to those of young stars. For the highest-luminosity ULXs,
the UVW1 fluxes may have an important contribution from the accretion disk. We
find that the majority of ULXs are associated with recent star-formation. Many
of the ULXs in our sample are located inside young OB associations or star-
forming regions (SFRs). Based on their colors, we estimated ages and masses
for star-forming regions located within 1 kpc from the ULXs in our sample. The
resolution of the OM was insufficient to detect young dense super-clusters,
but some of these star-forming regions are massive enough to contain such
clusters. Only three ULXs have no associated SFRs younger than $\sim$50 Myr.
The age and mass estimates for clusters were used to test runaway scenarios.
The data are in general compatible with stellar-mass binaries accreting at
super-Eddington rates and ejected by natal kicks. We also tested the
hypothesis that ULXs are sub-Eddington accreting IMBHs ejected by three-body
interactions, however this is not supported well by the data.

X-rays:binaries - ultraviolet: stars - galaxies: star clusters: general -
accretion, accretion disks

##  1 INTRODUCTION

Ultraluminous X-ray Sources - extremely bright X-ray sources with bolometric
luminosities exceeding the Eddington limit for a 20 M☉ object - continue to
puzzle astronomers. Due to their distance, located in external galaxies, it is
difficult to identify optical counterparts, even with the Hubble Space
Telescope (Ptak et al., 2006; Ramsey et al., 2006; Roberts et al., 2008; Tao
et al., 2011; Gladstone et al., 2013). Indeed very few objects have been
detected in the optical with high confidence $-$this is compared with the
hundreds detected in the X-rays111The following is a list of references for
some of the most famous objects with optical counterparts: NGC 5204 X-1 (Liu
et al., 2004), NGC 1313 X-1 (Yang et al., 2011), ESO 243-49 HLX-1 (Soria et
al., 2012) NGC 1313 X-2 (Zampieri et al., 2004; Liu et al., 2007; Ripamonti et
al., 2011; Zampieri et al., 2012), M81 X-6 (Liu et al., 2002; Swartz et al.,
2003; Moon et al., 2011), Holmberg IX X-1 (Grisé, Pakull, & Motch, 2006; Moon
et al., 2011; Grisé et al., 2011), Holmberg II X-1 (Kaaret et al., 2004; Tao
et al., 2012b), NGC 5408 X-1 (Lang et al., 2007; Grisé et al., 2012), M101 X-1
(Kuntz et al., 2005; Liu, 2009), NGC 4559 (Soria et al., 2005), two ULXs in
M51 (Terashima et al., 2006), NGC 2403 X-1 (Roberts et al., 2008), IC 342 X-1
(Feng & Kaaret, 2008), the ULX in NGC 247 (Tao et al., 2012a), NGC 6946 X-1
(Kaaret et al., 2010) and ULX P13 in NGC 7793 (Pakull et al., 2010; Motch et
al., 2011).. In addition, only three objects have established optical periods
that have been measured: M82 Kaaret et al. (2006); Kaaret & Feng (2007), for
NGC 1313 X-2 Liu et al. (2009); Zampieri et al. (2012) and NGC 5408 X-1
(Strohmayer, 2009; Han et al., 2012).

Multiple authors have shown that ULXs are associated with star formation in
their host galaxies (Ranalli, Comastri, & Setti, 2003; Grimm et al., 2003).
Indeed such a correlation has been observed for high-mass X-ray binaries
(HMXBs) in our own Galaxy (Grimm et al., 2003). Bodaghee et al. (2012) found
an average offset of $\sim$400 pc between galactic HMXBs and nearby SFRs (see
also Coleiro & Chaty, 2013). Many authors have found evidence that ULXs are
generally associated with large clusters and that the ULX is nearby these
clusters. For instance Kaaret et al. (2004) found that in starburst galaxies,
X-ray sources are, in general, located near star clusters but also that X-ray
sources with luminosities $>$ 1038.0 erg s-1 tend to be even closer to
clusters. Similar results were found by Rangelov et al. (2011) in the
starburst galaxy NGC 4449. Eleven X-ray binaries were found to be located
nearby or inside very young clusters. In the Antennae galaxy, 10 out of 14
ULXs seem to be associated with young stellar clusters (Zezas et al., 2002).
Clark et al. (2007) found 7 ULXs associated with clusters in the Antennae
galaxy using infrared images. These authors found that, in general, X-ray
sources tend to be close to large clusters (see also Clark et al., 2011).
Swartz et al. (2009) used photometric data from the Sloan Digital Sky Survey
to look for possible associations of 47 ULXs with SFRs or young superclusters.
They found that statistically ULXs are indeed associated with recent star
formation (within 100 pc distance), but no superclusters were detected given
the poor spatial resolution of the instrument. Poutanen et al. (2012)
performed spectral and photometric analyses of clusters associated with the
Antennae ULXs and found that almost all are very young (2.4 to 3.2 Myr), and
that only one resides inside a cluster (see also Rangelov et al., 2012).
Originally M82 X-1 was thought to be located inside a supercluster, until Voss
et al. (2011) showed that it is actually offset from the cluster. However, in
the same paper, the authors found that a ULX in NGC 7479 was associated with a
young supercluster, so such objects are known to exist.

In theory the location of a ULX in relation to its surrounding star clusters
can tell something about the environment in which the black hole was born as
well as constrain some of the properties of the black hole. For instance, if a
small black hole is born in a cluster of stars, the initial explosion can be
asymmetric enough to kick the black hole out of the cluster. This is known as
the runaway binary scenario (Zezas & Fabbiano, 2002). This theoretical
scenario does not work for larger black holes, such as intermediate black
holes (IMBH Colbert & Mushotzky, 1999), since a) the black hole is too big to
be susceptible to such kicks and b) even if such a black hole were kicked it
would return to the cluster on very short timescales because of gravitational
pull (Portegies Zwart et al., 1999, i.e. vkick $<$ vescape). However, there
are ways to kick intermediate black holes out of clusters using 3-body
interactions (e.g. Poutanen et al., 2012). In this scenario the intermediate
mass black hole and donor star can be kicked out of the cluster by another
young massive interloper star. Assuming this theory is correct, we would
expect to find most potential intermediate mass black holes inside young,
dense clusters of stars. Such, an environment would also readily explain the
growth of an intermediate mass black hole through stellar collisions with the
stars in the parent cluster (e.g. Gürkan et al., 2006). Thus, in summary, all
of this information can be combined to estimate characteristics of the ULX
including the age and mass of a stellar companion assuming the black hole
originated in a nearby cluster of stars or parent cluster (Zezas & Fabbiano,
2002; Kaaret et al., 2004).

The observational evidence for a possible association between ULXs and the
star-formation in the galaxy is still controversial. Some starburst galaxies
such as the Antennae galaxy, the Cartwheel galaxy and M82 contain an unusually
large number of very bright ULXs. On the other hand, there are many star-
forming galaxies without any known ULXs. In addition there are dwarf galaxies
(such as Holmberg II and Holmberg IX) and even some ellipticals that contain
bright ULXs. The ULXs found in elliptical galaxies seem to be fainter. Some of
the bright ULXs that appear may actually be interlopers (Irwin, Bregman &
Athey, 2004). Interestingly, most of the galaxies mentioned here (Antennae,
Cartwheel, M82, Holmberg II, and Holmberg IX) are either merging or
interacting with other galaxies. More recent examples are the colliding galaxy
pair NGC 2207/IC 2163 where 21 ULXs were detected (Mineo et al., 2013), Arp
147 with 9 ULXs (Rappaport et al., 2010), and NGC 922 with 12 ULXs (Prestwich
et al., 2012). The latter two are drop-through ring galaxies, similar to the
Carthwheel.

In this paper we use XMM-Newton archival images taken with the Optical Monitor
(OM) to explore ULX environments in nearby galaxies. Ultraviolet (UV) emission
is well suited to study star-forming regions (SFRs) and young clusters around
ULXs. The ULX sample presented here was extracted from the XMM sources
analyzed in Winter et al. (2006, hereafter, WMR). We selected all the ULXs
with unabsorbed luminosities LX $\geq$ 2.7$\times$1039 erg s-1, as estimated
by the authors. This limit corresponds to a 20 M⊙ black hole radiating at the
Eddington limit. Our ULX sample consists of all the sources for which OM UV
data was found in the archives. We added to this sample the ULX associated
with the MF16 nebula in NGC 6946 as a standard, since it is well-studied in
the literature. The original ULX sample in WMR was selected from galaxies
closer than 8 Mpc. The goal was first to examine whether ULXs are located
inside clusters or SFRs or are possibly related to nearby such regions. Second
was to use photometry to impose constraints on the nature of the optical
companions and the accretion mechanism of the ULX. We measured fluxes for 100
pc regions at ULX positions and for bright sources (SFRs) detected nearby.
Section 2 describes the data found in the archives and the photometry
procedure. We present the photometry results for the 100 pc region centered on
ULXs in Section 3. The results for the SFRs found close to ULXs are presented
in Section 4, together with the population synthesis modeling used to estimate
ages and masses. These are used in Section 5 to test runaway binary scenarios.
Finally, in Sections 6 and 7 we discuss our results and present the main
conclusions.

##  2 OBSERVATIONS AND DATA REDUCTION

OM is a 30-cm optical/UV instrument co-aligned with the X-ray instrument on
the XMM-Newton telescope. The OM includes three optical and three UV filters:
V, B, U, UVW1, UVM2, and UVW2, with effective wavelengths of 5430Å, 4500Å,
3440Å, 2910Å, 2310Å, and 2120Å, respectively. The OM instrument has a smaller
field of view than the X-ray instrument, but was not used by default in all
observations. We found 26 out of 36 ULXs from WMR that had been observed by
the OM and which met our criteria for selection (ULX is defined as any object
with LX $\geq$ 2.7$\times$1039 erg s-1, see Section 1 for details). We also
included the ULX in NGC 6946 as a reference object, because it is well studied
in the literature. All 27 objects are located in 16 nearby galaxies. We found
a total of 36 useful OM observations of these galaxies in the archives as of
2012 Jun 25 (meaning that some ULXs have more than one observation). These are
listed in Table 1. The ULXs were observed with a diverse set of filters. Most
ULXs were not observed with all of the filters, however. We discarded images
with artifacts like “ghost images”. However, with the exception of NGC 4736
XMM1 and NGC 5408 XMM1, all ULXs have images with the most sensitive of the UV
filters, UVW1. For the large majority of ULXs we found images in at least
three filters, allowing the construction of color-color plots. When redundant
images were available we used those images without artifacts and with longer
exposures.

We used the “sky images,” which are processed images created by the OM
pipeline. We used the XIMAGE software included in the HEAsoft v6.12 package to
measure the count rates in each filter. We converted the count rates into flux
values using the count rate to flux conversion factors listed in the XMM-
Newton Science Analysis System Users
Guide222http://xmm.esac.esa.int/external/xmm_user_support/documentation/sas_usg/USG/.
We applied the coincidence-loss and dead-time corrections on each pixel of the
extended source, following the procedure described in the XMM-Newton Users
Handbook333http://xmm.esac.esa.int/external/xmm_user_support/documentation/uhb/.

ULX 100 pc regions: To do the astrometry, all of the OM images were taken from
Kuntz et al. (2008), who derive coordinates for each image based on the
USNO-B1 catalog which is tied to the International Celestial Reference Frame
(ICRF). Kuntz et al. (2008) estimates the errors on these positions to be
0.3$\arcsec$. When an optical counterpart was observed, small shifts in the
images were made to align the OM and Chandra images further, thereby deriving
the absolute position of the ULX to less than the 0.3$\arcsec$ accuracy of the
Kuntz et al. (2008) images .

For each source, we created regions centered on the ULX position with
diameters of 100 pc, assuming the distances listed in Table 1 of WMR (two
examples are shown in Figure 1). The 100 pc regions translate into angular
sizes ranging from 2.6$\arcsec$ to 6.6$\arcsec$. The 100 pc regions are the
smallest possible regions that could be extracted for photometry since the
FWHM of the point spread function (PSF) for the OM instrument ranges from
1.5$\arcsec$ to 2.0$\arcsec$. The only exception to the 100 pc extraction
region was IC 342 XMM3 which is located inside a very UV-bright nuclear
cluster. For this ULX we used a region of 265 pc diameter instead of 100 pc,
in order to include the extended emission.

Surrounding Star Forming Regions: For each ULX we also measured fluxes from
nearby SFRs, using circular or elliptical apertures (of varying sizes, with
offsets up to 1 kpc from each ULX). The number of these regions, their sizes,
and offset distances vary for each object. For most ULXs we measured fluxes
for 3$-$5 nearby SFRs (see Figure 1 for two examples). In general, they are
large star-forming regions with diameters of up to 500 pc. The regions are
almost always brighter than the 100 pc regions around the ULX and the fluxes
measured are more accurate as a result. Aligning the sources between filters
is straight forward, since there are multiple stellar sources with which to do
the relative astrometry between filters.

Errors and Corrections: Most of the UV light from the ULXs is extended (i.e.
not a point source) and most of the emission is very faint. Because of this,
the resultant flux from the 100 pc regions could include multiple sources
rather than a single source that could easily be fit with the OM point source
extraction software. This leads to confusion in the measurement for the 100pc
region around the ULX.

Additional errors are introduced through aperture effects for both the ULX
100pc regions and the chosen star forming regions. These effects result from
using the same-sized apertures for each extraction region across all
wavebands. These color-based aperture corrections are complicated by the
extended emission (i.e. that the emission is not from a point source). Indeed
the needed corrections are higher for colors calculated between the optical
and the UV filters, since the point spread functions (PSFs) between these
filters are different.

In Figure 2 we calculate the theoretical color offsets obtained when using
apertures of the same sizes. We use the latest PSF calibration files available
from OM444ftp://xmm.esac.esa.int/pub/ccf/constituents/OM_PSF1DRB_0010.CCF. The
photometry for OM is calibrated for extraction apertures with radii of 12 and
35 pixels, for the optical and UV, respectively. The OM PSFs vary with source
brightness or, in other words, the ratio of the count rate to frame rate. The
FWHM decreases with source brightness, and therefore the offsets for the
plotted U - UVW1 color also decrease. The offsets between optical and UV
filters are similar, as illustrated by the B - UVW1 color in red (compared
with U - UVW1). The color offsets between the various optical filters are very
small because the PSFs are quite similar, however there is a slight offset
between the UVW1 and UVW2 filters, which is much smaller than the optical to
UV offsets, but which is large enough to warrant correction (blue lines).

Based on the theoretical results shown in Figure 2, we introduced color
corrections. However we note that, since practically all of the sources in
this paper are extended, the correction is approximate. We corrected the
measured colors between the optical and UV filters (U-UVW1 and B-UVW1) by
interpolating on the source flux up to $-$0.15 mag for faint sources (see also
Figure 2). For the UVW1 - UVW2 color, we used a maximum offset of 0.06 mag for
faint sources, scaling to zero for the bright sources. We found that these
corrections give consistent estimates for the cluster ages when other colors
are available to make the same calculation (e.g. M101 which has been observed
in all filters). In general these corrections are not large and are in fact
approximately the same size as the measurement errors on the original
magnitudes.

All extraction regions (both 100 pc ULX regions and SFRs) were background
corrected. The background was measured by extracting fluxes from nearby areas
that were free of visible point sources.

The color excess values E(B-V) for the Milky Way extinction were obtained from
the NASA Extragalactic Database (NED). Both Galactic and extragalactic
extinctions were calculated using the extinction curves from Cardelli,
Clayton, & Mathis (1989), assuming R${}_{V}=$ 3.1. All photometric magnitudes
presented were corrected using Galactic extinctions and the intrinsic
extinctions obtained from the color-color models presented in Section 4.

##  3 PHOTOMETRY RESULTS FOR THE 100 PC REGIONS

In this Section we use the fluxes measured from the 100 pc regions around the
ULX to check for possible UV counterparts. These could be single stars (the
companions in the ULX binaries), small OB associations or larger star-forming
regions (clusters). An important contribution to the optical or UV flux could
also come from the accretion disk. We first use the UVW1 fluxes to make
general estimates for the different possible counterparts. Then, where
available, we use the colors in different OM bands to compare with models of
standard stars and accretion disks.

###  3.1 UVW1 Magnitudes

The UVW1 is the most sensitive of the UV filters and most of our ULXs have
images taken with this filter. NGC 4736 XMM1 and NGC 5408 XMM1 have no images
in the UVW1 band. For these ULXs we used the U and UVM2 band, respectively.
For ULXs with multiple observations we used the error-weighted average. Based
on the XMM-Newton Users Handbook, we estimate the detection limit for a O5 V
star to be $\sim$18.8 mag in the UVW1 filter. This means an individual O5 V
star will be detected to about 2.4 Mpc with the UVW1 filter if the exposure
time is 1000 s. Our ability to measure single stars around ULXs is difficult
because the brightness of these stars is very close to the detection limit of
the UVW1 filter, though the average exposure time is longer (3000 s). The
magnitudes provided in Table The Local Environment of Ultraluminous X-ray
Sources Viewed by XMM Newton’s Optical Monitor should therefore be treated as
upper limits. The measured UVW1 magnitudes are later used to estimate the
nature of the UV emission (i.e. whether the flux is consistent with values
expected from a single star, multiple stars, or accretion emission).

####  3.1.1 Single stars

In Column 11 of Table The Local Environment of Ultraluminous X-ray Sources
Viewed by XMM Newton’s Optical Monitor, we list the estimated stellar type
based on the UVW1 magnitudes, thereby assuming that all the flux within the
100 pc region is emitted by a single star. Given the signal-to-noise (S/N)
ratios in Column 9 of Table The Local Environment of Ultraluminous X-ray
Sources Viewed by XMM Newton’s Optical Monitor (we take S/N $>$ 3 for a
detection), we estimate that 7 ULXs may have single stars as counterparts. An
additional two ULXs have no significant detections. For these we show
3$\sigma$ upper limits. In Column 12 of Table The Local Environment of
Ultraluminous X-ray Sources Viewed by XMM Newton’s Optical Monitor we also
show the stellar types for all previously identified optical counterparts in
the literature. We note that while the absolute magnitude is consistent with
that of a single star, the UV emission from these objects are also likely
contaminated by accretion disk emission as has been shown in many cases
(Roberts et al., 2011; van Paradijs & McClintock, 1994).

Ten of our objects have estimated spectral types for donor stars in the
literature. Of these four ULXs (NGC 1313 XMM3, NGC 2403 XMM1, Holmberg IX XMM1
and M101 XMM3) have UVW1 spectral-type predictions that are identical or very
close to that found in the literature. For the other six ULXs which have
previous optical counterparts in the literature, the 100 pc UV fluxes suggests
multiple stars. In these cases the ULXs are located inside or very close to
small clusters or OB associations and therefore the UV fluxes are likely
contaminated by surrounding star formation. The following outlines background
information for ULXs which may suffer contamination:

IC 342 XMM1: IC 342 XMM1 has a F8-G0 Ib supergiant as an optical counterpart
according to Feng & Kaaret (2008). They find no O stars within 300 pc of the
ULX, and estimate that for detected stars, the minimum stellar age is 10 Myr.
If this star is indeed a G0, then the absolute magnitude of the star is only
$-$3.1, and is well below the detection threshold of the UVW1. We estimate
that the UVW1 flux is equivalent to five O5V stars, but given the previously
estimated age, the UV emission is likely produced by several later type stars.

Holmberg II XMM1: Homberg II XMM1 (HoII X-1) has a 21.5 V-band magnitude
stellar counterpart with colors consistent with O stars on the main sequence
or B supergiants (Kaaret et al., 2004). The ULX is located inside an
association of young stars (Kaaret et al., 2004). We suspect that here too,
the total UVW1 flux is overestimated due to contamination.

NGC 5204 XMM1: NGC 5204 XMM1 (better known as X-1) has a B0 Ib supergiant
counterpart, based on the Si III $\lambda$1299 line (Liu et al., 2004). The
Hubble images in Liu et al. (2004) show at least 3 star clusters and OB
associations within our 100 pc region (4.3$\arcsec$). We estimate
approximately eight O stars from our UVW1 flux, though here too the flux is
likely contaminated.

NGC 5408 XMM1: NGC 5408 XMM1 is also located very close to an OB association.
Grisé et al. (2012) estimated its age at 5 Myr and also identified a B0
supergiant as the likely ULX counterpart. Several young stars from the OB
association fall into our 100 pc extraction region. Instead of a single star,
we find the UVW1 flux is equivalent to 8 O5V stars.

NGC 247 XMM1: Tao et al. (2012a) used HST observations to show that the ULX is
very close to a large and relatively young stellar association. The stars
closest to the ULX are contaminating the UVW1 flux. The estimate is therefore
equivalent to 5 O5V stars.

M51 XMM6: The ULX is located at the edge of a large SFR, its emission
dominates our 100 pc flux. The measured UVW1 flux is equivalent to 35 O5V
stars.

NGC 6946 XMM1: This ULX is associated with the bubble nebula MF16 and a UV
point source was identified using HST data by Kaaret et al. (2010). Several
nearby bright stars and the nebular emission are contaminating our UVW1 flux
here as well (see also Berghea & Dudik, 2012). The measured UVW1 flux is
equivalent to $\sim$ 3 O5V stars.

Finally, for three ULXs (not including the upper limits) we have no previous
published information. For these, even if our estimates are compatible with
single stars, the UV images suggest extended star-forming regions, even beyond
the 100 pc regions. The UVW1 emission here is more likely produced by many
late-type stars, instead of a single young star, given the spatial information
we have.

####  3.1.2 Multiple stars or star-forming regions

For 18 ULXs we measured UVW1 fluxes that are too high for single stars, they
seem to be located inside small OB associations or SFRs. For these, Column 11
of Table The Local Environment of Ultraluminous X-ray Sources Viewed by XMM
Newton’s Optical Monitor lists the equivalent number of O5V stars, assuming
that all the UVW1 flux inside the 100 pc regions comes from such stars. Two of
the most obvious examples are IC 342 XMM3 and NGC 4736 XMM1.

IC 342 XMM3: IC 342 XMM3 is very close to the nucleus of the galaxy. Kong
(2003) estimated that it is located 3$\arcsec$ from the dynamical center of
the galaxy. Böker, van der Marel, & Vacca (1999) measured the nuclear cluster
in IC 342 using infrared and optical observations and found a mass $<$
6$\times$106 M☉, and a relatively young age (106.8-7.8 yr), but found no
evidence for a supermassive black hole. If this age estimate is correct,
O-stars are too young to be found in the region. Indeed, using the models of
Leitherer et al. (1999) we find O-stars in the number suggested will produce
the emission observed in less than 1 Myr. By 106.8-7.8 years (the lower limit
from Böker, van der Marel, & Vacca, 1999), all of these O-stars would be dead.

NGC 4736 XMM1: For NGC 4736 XMM1 we estimate 747 O5V stars. The ULX is located
inside the face-on spiral galaxy NGC 4736. HST images show strong unresolved
star formation at the ULX location, and our estimated extinction is also high
(Table 4).

####  3.1.3 Contributions from accretion disks

In Fig. 3 we plot the UVW1 magnitudes against the bolometric luminosities for
all of the ULXs in our sample (bolometric luminosities were taken from WMR).
The diagonal line in the upper-right corner of this image represents the
expected accretion disk emission in the UV as a function of the bolometric
magnitude. This is a theoretical estimate, calculated using a standard disk
model for a ULX accreting at the Eddington limit (Shakura & Sunyaev, 1973).
The dotted line represents the emission from an irradiated accretion disk,
which is also the maximum contribution to the UV emission we would expect from
a disk (Sanbuichi et al., 1993; Berghea et al., 2010a).

The position of this line implies that a significant portion of the measured
UV flux is likely from accretion disk emission if the ULX has a bolometric
luminosity in excess of 1040.0 erg s-1. This is also likely the case for
sources appearing close to the theoretical model in the figure, such as NGC
1313 XMM3 and Holmberg IX XMM1 (see Fig. 3). Interestingly these results are
consistent with the findings of Roberts et al. (2011), who suggest, based on
optical spectroscopy, that some portion of the UV emission from NGC 1313 XMM3
and Holmberg IX XMM1 comes from a disk component. We note that these ULXs also
show some of the highest X/UVW1 ratios in Column 10 of Table The Local
Environment of Ultraluminous X-ray Sources Viewed by XMM Newton’s Optical
Monitor.

An even larger fraction of the UV emission would be expected to come from the
disk if it were irradiated. In fact, the UV emission can be increased by up to
a factor of 10 for stellar mass BHs as shown in the Figure (Berghea et al.,
2010a). If this is the case for NGC 1313 XMM3 and Holberg IX XMM1, it would
imply that the UVW1 flux is almost entirely from the disk.

We also note that, unlike typical AGN, we do not see a trend in X/UVW1 (where
X is a proxy for the bolometric luminosity). Indeed there is no correlation at
all for either single star sources or multi-star sources.

While we cannot rule out a strictly stellar origin for the ULXs in our sample,
it is very likely that a fraction of the UV emission we observe is from the
accretion disk. However as Fig. 3 shows, for most of our sources, the
accretion disk is not bright enough to be responsible for all of the UV
emission (even if it is irradiated). As mentioned previously (and as Table The
Local Environment of Ultraluminous X-ray Sources Viewed by XMM Newton’s
Optical Monitor, column 11 shows) most of our sources are bright enough that
the emission is likely to be contaminated by multiple stellar sources within
the 100 pc aperture, thereby making it impossible to separate the two. In
addition, for those sources with stellar estimates, the UV emission is bright
enough to come from both the donor star and a disk. Since one of the main
goals of this paper is to qualify the environments around the ULX, we feel
that the approximate spectral type is very important to understanding the
origin and age of the environment from which the ULX was born. Again, this
does not mean that some of the total emission does not come from the accretion
disk, but that the photometry looks like a specific type of star.

####  3.1.4 UVW1 variability

Some of the objects in this sample have more than one observation taken in the
UVW1 band, however only 4 objects have more than one observation in UVW1 where
the source was detected (many of our objects have multiple observations but
the others are non-detections). We therefore looked into the variability of
these four objects. The results of the analysis are shown in Table 3. Of the
four objects, three (IC 342 XMM1, M 101 XMM2, and NGC 1313 XMM3) have only two
observations each in UVW1. As can be seen in Table 3, the two magnitudes
observed are the same to within the errors for all three sources. These
objects can therefore be variable only to within this error based on the
limited observations available here. The fourth source however, Holmberg II
XMM1, is the only source with more than two observations in UVW1. In this
case, the magnitude appears to vary by 0.25 mag. over 8 years. A $\chi^{2}$
test suggests that this source may indeed be variable ($\chi^{2}$ = 102 for 3
degrees of freedom, assuming a constant model).

###  3.2 OM Colors

We performed a color comparison of the 100 pc regions with standard stars. Six
color-color plots are presented in Figures 4-6. To create the stellar models,
we used the Kurucz star library555http://www.stsci.edu/ftp/cdbs/grid/k93models
to estimate OM colors for stars with solar metallicity. The fluxes were folded
through the latest OM response matrices. Kuntz et al. (2008) generated an OM
Catalog and found that in general the OM colors match the Kurucz models well,
except for late type stars, where a systematic offset is seen. We note that
not all objects were observed in every filter. Therefore some ULXs appear in
some plots and not in others. If a ULX appears in one plot and not another,
this simply means that the ULX was not observed in one or more of the color
filters needed for the second plot.

Since the accretion disk can have an important contribution to the color
emission, we show the theoretical estimates from a standard accretion disk.
The UV emission in this case is described by a power-law with energy index 1/3
for any BH mass. Therefore the UV colors do not change with the BH mass. They
change however when we include self-irradiation of the accretion disk. We
followed Sanbuichi et al. (1993) to calculate theoretical spectra from
irradiated disks for BH masses up to 10000 M⊙, and then estimated UV colors.
We assumed accretion at the Eddington limit. The theoretical OM color tracks
are plotted as thick orange lines in Figures 4-6. The errors in the figures
are large, especially for the 7 ULXs where the 100 pc flux is compatible with
single stars (not including the upper limits in Table The Local Environment of
Ultraluminous X-ray Sources Viewed by XMM Newton’s Optical Monitor). The color
information was sufficient to determine spectral type for only 2 sources:
NGC4395 XMM1 and M101 XMM2. We find a spectral type of B3I and O3 - O5V
respectively. Both estimates are close to the UVW1 predictions, but do not
completely match them. We note that in many cases, our fluxes are likely not
from single stars, as most ULXs are located in crowded fields (see Section 3.1
above). We also note that the self-irradiated disk produces colors that are
similar to young stars. It is possible that a significant part of the detected
UV flux for some ULXs plotted in Figs. 4 through 6 is emission from the
accretion disk. This is consistent with the accretion disk contamination we
discussed in Section 3.1.3.

We note that the irradiation of the donor by the X-ray emission further
complicates the optical photometric estimates. Copperwheat et al. (2005) found
that both the absolute optical magnitudes and the colors of ULXs (accretion
disk plus companion) vary significantly with the mass of the BH. For example,
for an O5V star companion, when the BH mass is varied between 10 and 1000 M⊙,
the V-band changes by 1.0 magnitude and the B-V color by 0.2.

##  4 STARBURST99 MODELING AND COLORS

In the previous Section we found that some ULXs are apparently located inside
SFRs. In most cases, bright SFRs are also found in the vicinity of ULXs. Such
regions usually contain young clusters or OB associations, which are strong UV
sources. Indeed these environments seem necessary for a ULX to subsist. For
example, IMBHs can form inside dense young clusters (Portegies Zwart &
McMillan, 2002; Gürkan et al., 2004). In addition, stellar-mass black holes
can produce super-Eddington luminosities only if the donor is a massive star
(Podsiadlowski et al., 2003). For this reason we explored any possible
correlations between the ages and masses of our star forming regions and the
ULX.

We estimated ages and masses for the nearby star-forming regions, using the
Starburst99 code (Leitherer et al., 1999) version 5.1. We simulated an
instantaneous burst model with a Salpeter Initial Mass Function (IMF) and
having exponent $\alpha=$ 2.35, mass boundaries at 1 and 100 M☉ and a total
cluster mass of 106M☉.

The ULXs in this sample are located in different galaxies having a range of
star formation rates (see WMR). However, Winter et al. (2007) used the
absorption features in XMM-Newton X-ray spectra to measure the abundances for
the ULXs in the WMR objects. The results are consistent with solar abundances
for the entire sample, despite the wide range of galaxy host properties. On
the other hand, some authors found lower abundances for some ULXs (e.g.
Brassington et al., 2005; Mapelli et al., 2009; Berghea et al., 2010b; Pintore
& Zampieri, 2012). Fortunately the models for young UV-bright stars are not
sensitive to metalicity. We therefore use solar abundances and do not expect
large errors from this assumption. Moreover, given the large error bars for
the measured colors, we found that models with different IMF slopes or cutoff
masses would not change our results.

###  4.1 Starburst99 Modeling

Apart from the 100 pc regions, we identified several bright SFRs close to most
of our ULXs. The total number of regions for all 27 ULXs is 90. We used the
generated Starburst99 models to predict ages and masses for both the 100 pc
and the star forming regions.

In Figures 7$-$9 we compare the OM colors of the SFRs with the Starburst99
models evolved to 900 Myr. The Starburst99 model tracks are also plotted for
predicted extinctions of A${}_{V}=$ 0, 1 and 2. Because most objects were not
observed in every filter, the ULXs and associated SFRs are not the same from
figure to figure. With the exception of Holmberg II and NGC 5408, all the
galaxies were observed with at least three different filters. We chose the
colors that make up each axis in such a way to ensure all regions are plotted
in at least one plot. The measured colors were corrected for aperture effects
as explained in Section 2. We show the magnitude and direction of the maximum
aperture correction applied to each plot with an arrow.

We estimated approximate ages by comparing the colors to the Starburst99
tracks. When colors for a region were present in more than one plot, we used
all the available data to make a best estimate for the age. Three of the six
color-color plots are degenerate with respect to the cluster age and
extinction (Fig. 8, both panels and Fig. 9, right panel). Therefore, when one
color-color plot was degenerate, we used information from the other plots to
establish a cluster age. We found that the estimated ages agree reasonably
well for different colors, to within the errors.

The mass of the cluster can be estimated from the age and extinction predicted
from Figures 7$-$9. To do this we first modeled the relationship between
cluster age and UVW1 absolute magnitude for a 106 M☉ cluster having three
extinction values (A${}_{V}=$ 0, 1 and 2). This model is shown in Figure 10.
We then find the UVW1 magnitude in this plot that corresponding to a cluster
age and AV from the Figures 7$-$9 in the analysis described above. The
magnitude from this plot will correspond to a mass of 106 M☉. We finally
scaled this mass based on our observed UVW1 magnitudes. For example, a 50 Myr
cluster with AV = 0 will have a UVW1 magnitude of -14 (See Figure 10). If our
observed UVW1 value is a factor of 10 fainter then we can predict that the
mass is smaller than 106 M☉ by a factor of 10. For the clusters with no
reliable age estimate, we calculated an approximate mass, assuming a fixed
cluster age of 50 Myr.

###  4.2 Results from the Color Comparison

The age and mass estimates from the color-color plots are listed in Table 4.
Columns 3 and 4 show the average extinction and age estimates of all
surrounding star forming regions. Only extinction values from the non-
degenerate color-color plots were used.

In Column 2 of Table 4 we show extinction values calculated using the column
densities obtained by WMR from spectral fitting. We assume E(B$-$V) $=$ 1 for
a column of 5.5$\times$1021 cm-2. By comparing the estimates from the X-rays
with the OM color estimates in columns 3, 4 and 8, it is evident that in
general the X-ray estimates are much higher. The average absorption column
density based on X-ray spectral analysis in WMR is 4.5$\times$1021 cm-2. This
corresponds to an extinction of A${}_{V}=$ 2.5, while most of our estimates
are $\leq$1\. Our results suggest that the absorption is localized, in the
immediate vicinity of the ULX (on scales less than the binary separation).
Krolik (2004) proposed that some ULXs could be IMBHs inside dense molecular
clouds and are fueled through Bondi accretion. For the most part, our color-
color plots do not show strong extinction usually associated with dense
interstellar media, and therefore do not support this scenario.

In Columns 6-9 we list the results for the 100 pc regions. In many cases,
reliable ages could not be estimated because of low fluxes and large errors.
However, the ones that could be measured indicate relatively young ages. Only
NGC 253 XMM1 showed an age of $>$100 Myr.

Columns 10-13 of Table 4 show the results for the closest SFRs to each ULX.
Thirteen (approximately half) are actually located inside these star forming
regions, meaning that the star forming region overlaps with the 100 pc ULX
regions. This is true even if the intrinsic offset is quite large (e.g. $>$200
pc for NGC 253 XMM6). These 13 objects are marked in bold in column 10. Four
of the ULXs have age estimates $<$10 Myr. Of these, the SFR located near NGC
4631 XMM1 and IC 342 XMM3 have relatively large masses ($>$105 M☉). Apart from
the ULXs located inside the SFRs, nine other ULXs have SFRs within 200 pc (21
in total out of 27).

The results for the youngest SFRs are shown in Columns 14-17. Seventeen out of
21 ULXs have nearby regions younger than 10 Myr. Only NGC 253 XMM1 does not
have a nearby SFR younger than 50 Myr.

Finally, we consider the most massive SFRs found nearby the ULXs in our sample
(cols. 18-21). There are 16 SFRs with masses $>$105 M☉, but only 12 with age
estimates. Of these 12, roughly half (7) are younger than 50 Myr, including
the central cluster in IC 342.

The final row of Table 4 shows the median extinctions, ages, offsets and
masses for the five categories described above: the average region, the 100 pc
region, the closest, youngest, and most massive star forming regions. We find
that in general the closest star forming regions have stellar ages that are
approximately 9 Myr and are offset from the ULX at an average distance of 87
pc. On the other hand, we find that the majority of the most massive regions
are 12 Myr in age and sit 240 pc from the ULX. The youngest star formation
regions tend to be 309 pc away and be approximatly 5 Myr, with the exception
of NGC 253 X-1 which has a relatively old population of stars.

###  4.3 Previously Known Star-forming Regions

Some of the SFRs located nearby the ULXs in our sample were previously studied
in the literature. These are discussed in detail below.

Holmberg IX XMM1: Holmberg IX XMM1 (known as X-1) is associated with a bubble-
like shell, and is located at the edge of a stellar cluster with estimated
mass of 103 M☉ and older than 20 Myr (Grisé, Pakull, & Motch, 2006). Grisé et
al. (2011) used Hubble data to estimate a color excess of E(B-V)$=$ 0.26. They
found that the ULX is located inside an OB association with mostly young stars
($<$20 Myr) and having mass $\sim$2$\times$103 M☉. We could not make a
reliable age estimate based on the OM colors due to insufficient data, but our
mass estimate (6$\times$103 M☉) is consistent with this result, given the
older age we used (50 Myr).

NGC 1313 XMM3: NGC 1313 XMM3 is located at the edge of a relatively young
cluster, which we estimated to be $\sim$5 Myr old (for extinction
A${}_{V}\approx$ 0.5, col. 15 in Table 4). Liu et al. (2007) found a similar
result: age $\sim$10 Myr for a slightly larger extinction (A${}_{V}\approx$
1). Tao et al. (2012a) used Hubble data to estimate the age of the stellar
association located very close to the ULX. Using a color excess of
E(B-V)$=$0.18 they found a mixed population of mostly young stars, in
agreement with our estimate.

IC 342 XMM3: For the central cluster where IC 342 XMM3 is located, we estimate
an age of 5 Myr, and an extinction A${}_{V}=$ 2.5. This age is slightly
younger than the one obtained by Böker, van der Marel, & Vacca (1999), and our
estimated mass for the cluster, 107 M☉, is almost double. However, the
distance to IC 342 used by these authors was smaller than our distance (3.9
Mpc), and this could account for the discrepancy. For consistency our
distances are taken from WMR.

NGC 5408 XMM1 Grisé et al. (2012) analyzed Hubble data for the OB association
located close to NGC 5408 XMM1. For an extinction A${}_{V}=$ 0.25 they
estimated an age of 5 Myr and suggest that the ULX could be a runnaway binary
with a velocity of $\sim$ 25$-$50 km s-1. Unfortunately we only have images in
two filters for this ULX, and could not get estimates of age for this object
using OM.

##  5 RUNAWAY SCENARIOS

One of the main findings of this study is that many of the ULXs in our sample
are located inside star forming regions rather than outside of them. We note,
however, that the OM has poorer resolution than the instruments used in many
of these studies. The rest of the ULXs (12 total), are indeed located nearby
such regions. In order to explain these rogue ULXs we explore the scenario
where ULXs form in young stellar environments and are later ejected.

###  5.1 Stellar-mass X-ray Binary Runaway

If the stellar mass black holes have super-Eddington accretion as described by
Begelman (2002), then Podsiadlowski et al. (2003) shows that the observed
X-ray luminosities can be obtained in systems with massive donors ($\gtrsim$
10 M☉, or stars earlier than $\sim$B2). For this analysis we assume that the
ULXs are stellar-mass X-ray binaries.

ULX stellar-mass BH candidates can be found outside of the parent cluster if
the birth of the black hole creates kicks of sufficient strength. Sepinsky et
al. (2005) estimated that $\sim$70% of the X-ray binaries can be ejected from
the cluster in this way. They also find that the younger the region ($\sim$10
Myr), the more X-ray binaries are present outside the cluster. For such
clusters, typical offsets are 10 pc$-$100 pc. Larger offsets (of $\sim$1000
pc) are found when the cluster is older ($>$50 Myr), but the number of
binaries is also much reduced in that case. Recently, Repetto et al. (2012)
found that natal kicks similar to neutron stars can explain the distribution
of black hole binaries in our Galaxy. However, Reid et al. (2011) and Miller-
Jones et al. (2009) use radio interferometry to measure the parallax of
Galactic black holes and find that relatively small kicks are required to
produce the velocities measured.  Zuo & Li (2010) simulated natal kicks for
X-ray binaries using different models of clusters and a range of parameters.
They found that the offsets are in general consistent with the observational
data from Kaaret et al. (2004).

####  5.1.1 Estimating kick velocities

To estimate the kick velocity, we followed the procedures of Fryer & Kalogera
(2001), who used Galactic neutron star velocities to scale to black holes
velocities. We note that Fryer & Kalogera (2001) assume that black holes are
formed after a neutron star is formed (i.e. not through direct collapse),
which may not be the case if these ULXs are more massive black holes.

The neutron star velocities we used when we performed this calculation are an
average of three more recent measurements by Cordes & Chernoff (1998),
Arzoumanian et al. (2002) and Hobbs et al. (2005). Cordes & Chernoff (1998)
found a velocity distribution consisting of a two-component Maxwellian
distribution with characteristic velocities of 175 km s-1 and 700 km s-1, and
a weight predominantly on the lower velocity component. The mean two-
dimensional and three-dimensional velocities for this model are 331 km s-1 and
421 km s-1 respectively. More recently, Arzoumanian et al. (2002) found a
similar model but with characteristic velocities of 90 km s-1 and 500 km s-1
and roughly equal weights. This model implies more pulsars with higher
velocities than the previous model. The most recent study of Hobbs et al.
(2005) fit the pulsar velocity distribution to a single Maxwellian
distribution with a mean three-dimensional velocity of 400 km s-1 and no
evidence for a multiple-component velocity. The mean pulsar velocity we use
for our calculations is 360 km s-1.

Using the Fryer & Kalogera (2001) method, we calculate kick velocities as:

$v=500($M${}_{BH}+$M${}_{comp})^{-1}$ km s-1.

The velocities are scaled to the total mass of the binary BH (MBH) plus the
companion (Mcomp). These velocities are more than twice the values found by
Zezas & Fabbiano (2002) who use a similar method for the Antennae ULXs. The
discrepancy likely results because Zezas & Fabbiano (2002) use the values of
Cordes & Chernoff (1998) for the initial neutron star velocity, while, we use
these values averaged with those of Arzoumanian et al. (2002) and Hobbs et al.
(2005). Our calculation suggests that binary BH/companion star systems could
have kick velocities that are quite large. Indeed Arzoumanian et al. (2002)
estimate that 50% of the pulsars have velocities greater than 500 km s-1.

We note that the velocities we can measure for our ULXs (see below) are two-
dimensional. This means configurations of ULXs that have been kicked from the
cluster toward or away from the viewer will not be taken into consideration in
this estimate. As a result the average velocity over all configurations of
offsets is 4/$\pi$ times lower than the three-dimensional values $-$ at most a
27% error on the velocity. On the other hand the neutron star velocities from
which we base this calculation have a deviation well above this 27% error. We
therefore do not correct for the effect in our calculations.

####  5.1.2 Estimating X-ray binary lifetimes

We estimate the lifetime of X-ray binaries using two mass transfer types (e.g.
Patruno & Zampieri, 2008). In the first case (A) we assume that the mass
transfer starts early when the donor star is on the main sequence, while in
the second (B), it begins after the donor has evolved off the main sequence.

It is much harder to observe a ULX in case B because the brighter/super-
Eddington phase is much shorter (by a factor in the range 10$-$100, see
Podsiadlowski et al., 2003; Rappaport et al., 2005). On the other hand, the
luminosities reached for case B are larger by an order of magnitude, thereby
making the likelihood of seeing the ULX higher since the selection of possible
companions is not restricted to massive stars. The lifetime for the ULX in
case B is equal to the lifetime of the donor star, as in Zezas & Fabbiano
(2002).

The lifetime calculation for case A requires a detailed stellar evolution
simulation (e.g. Podsiadlowski et al., 2003), which is beyond the scope of
this paper. A simple estimate of the lower limit can be obtained as follows.
We assume that the ULX luminosity is super-Eddington by a factor of 10 and
that the accretion efficiency, $\eta$, is approximately equal to 0.1. If the
donor has an initial mass, M, and a luminosity, L (ten times the Eddington
limit), we estimate the lifetime for this system by taking the harmonic mean
between the lifetime of the donor on the main sequence and the time T
$=\eta$Mc2/L, in which the entire mass of the stellar companion is accreted.
These two mass transfer cases are taken as limiting cases. The transfer can
theoretically start at arbitrary times on the main sequence. However, Poutanen
et al. (2012) point out that there is a minimum delay between the time the
binary is ejected and the age of the donor, because the primary evolves inside
the cluster before the supernova explosion. The binary must therefore be at
least 3 Myr old when the system is ejected. We take this into account when we
calculate the offsets by subtracting 3 Myr from the donor’s remaining
lifetime.

####  5.1.3 Estimating theoretical offsets

Using the velocities and lifetimes calculated in this way we can estimate the
maximum offsets from the star-forming regions for different BH and donor
masses, and compare them with the measured offsets for our sample. These
calculations are bound by two constraints: the age of the companion and the
accretion rate of the system. The ages of the companion stars are constrained
from the Starburst99 models and observed OM data of the cluster from which
they are thought to be born. This in turn constrains the offset distances for
the ULXs, assuming the stellar companion was acquired from the closest
cluster. The accretion rates are estimated using the bolometric luminosities
from WMR.

1\. Age estimates for the clusters: Based on the derived age (presented in
Table 4), we plot in the left panel of Figure 11 the largest mass possible for
the companion, assuming it was formed in this cluster. All objects with age
estimates are plotted here. When several star-forming regions were identified,
we used the data for the closest to the ULX, i.e. those presented in cols.
10-13 of Table 4. For the case A transfer, we estimate offsets based on the
lifetime of the binary and assume super-Eddington luminosities for the
accretion rate. For case B the lifetime is the natural lifetime of the star.
In case B the accretion rate is negligible because it only comes into play at
the final phase of the campanion star lifetime (consists of $\sim$100,000
years). Only 9 ULXs could have companions with masses $\gtrsim$ 10 M☉ (earlier
than type B1), capable of sustaining super-Eddington luminosities on long
timescales (case A mass transfer).

In Figure 11 we compare the measured (projected) offsets with the theoretical
offsets derived from the upper-limit calculation described in this section for
different BH masses. We note that the theoretical lines presented in this plot
do not include the potential of the cluster. We estimate that the cluster
contribution to the total velocity will be very small for large kick
velocities ($>$ 40 km s-1). However for small kick velocities this
contribution could be quite large and if the kick velocity is small enough the
BH will not escape the cluster at all.  In the case A transfer the offsets
plotted are consistent with BH masses $<$20 M☉. For the case B transfer,
larger offsets are predicted by the theoretical models, and most are
consistent with BH masses $<$40 M☉, which we take as the mass limit for a BH
formed from a supernova. We also show the kick velocities for each donor mass
at the bottom of the plot (in km s-1).

As an example consider NGC 247 XMM-1. For the closest star cluster region to
this ULX (64 pc offset from the ULX position) we find a 12.5 M☉ donor that is
10 Myr old (see the results for this object in Table 4). This places the ULX
firmly between the green and red dotted lines in Figure 11. The lines
represent the expected relationship for 20 M☉ and 10 M☉ black hole systems
respectively based on our models. This implies that for Case A mass transfer
(dotted lines), the accretor has a mass of $\lesssim$15 M☉. The ULX donor mass
vs. offset is well below any of the case B mass transfer (solid lines) models,
and therefore we have no constraint on the mass for this case.

In general our measurements follow the relationship between the mass of the
companion and the measured offset distance as suggested by the theoretical
offset tracks for the case B (continuous lines). We note that the errors in
estimating the masses of the objects used for Figure 11 are very large. As
mentioned in Section 2, the resolution of the OM is not sufficient to resolve
the ULX and separate the ULX emission from the surrounding stars. Therefore
the colors used to estimate the stellar mass are likely contaminated. Moreover
these plots only show the closest star forming regions to the ULX, when in
fact, the ULXs could be related to other regions farther away.

2\. Using the bolometric luminosities from WMR (the right panel of Figure 11):
This calculation does not require an estimate for the cluster or companion
star age. The offset calculated is based solely on the timescale of the
accretion for the system. The accretion rate is estimated from the bolometric
luminosities of WMR. Case B is not relevant here because the accretion time is
such a small part of the lifetime of the star (100 Kys vs. 10 Myr). We compare
the empirical offsets of the closest regions for all ULXs with those
calculated (theoretical) using the bolometric luminosity method described
here. The theoretical tracks were calculated for different donor masses, and
two different accretion rates: one at the Eddington luminosity (dashed lines)
and one at ten times this limit (continuous). The plot is similar to Fig. 3 in
Kaaret et al. (2004), but differs because these authors assumed the same kick
velocity (10 km s-1) for any BH mass. For low-velocity kicks, the ULX might
not be able to escape if it formed in a dense cluster. We mark the location of
a kick velocity of 10 km s-1 in the plot by thick black lines crossing the
theoretical tracks. ULXs located on the right side of these tracks have a low
probability to escape such a cluster.

In general, the ULXs follow the theoretical predictions, showing smaller
offsets for larger luminosities. The super-Eddington tracks (continuous
lines), as upper limits, permit larger offsets for the same luminosity. This
is because the kick velocity is higher for smaller black holes (see the top
X-axis for the offset relationship to BH mass). A comparison of the donor mass
plot with the bolometric luminosity plot suggests that the latter calculation
predicts smaller offsets. This might be because the donor mass plots include
models for Case B, where the luminosity models do not. Most of the ULXs
plotted are consistent with the super-Eddington case, but not with accretion
at the Eddington limit. There are a few ULXs with high luminosities that show
larger offsets than permitted even for the super-Eddington accretion.

NGC 247 XMM1 has bolometric luminosity of 7.1$\times$1039 erg s-1. For sub-
Eddington accretion, the corresponding mass of the BH is $>$55 M☉, and the
offset is too large for any companion mass. If the accretion is super-
Eddington, the companion could be a massive star ($\sim$20 M☉).

###  5.2 IMBH Runaway

If ULXs contain IMBHs then they cannot be ejected from dense clusters by natal
kicks alone, because the kick velocity is smaller than the required ejection
velocity. Miller & Colbert (2004) presented other mechanisms that can, in
principle, deliver higher kicks. The authors discuss the scenario of three-
body interactions in young clusters, and find that significant kicks ($>$10 km
s-1) can be achieved even for a large BH (1000 M☉), if the interloper is a
young massive star (an O star). The kick velocity will depend also on the
binary orbital velocity, which can be estimated assuming Roche lobe overflow.

Poutanen et al. (2012) also suggested three-body interactions to explain the
association between the Antennae ULXs and very young clusters. The authors
argue that the natal kicks cannot explain the data because the binary cannot
be ejected before at least 3 Myr (the life of a massive star). This leaves no
time for the binary to reach the present position outside the cluster. Mapelli
et al. (2011) performed simulations of three-body interactions in massive
clusters. They investigated ejections of black hole binaries with large
masses, capable of achieving the high accretion rates required to explain ULXs
(up to 80 M☉). They found that almost half of these binaries are expelled
within 10 Myr, and that the offsets are consistent with ULX observational
data.

Following Miller & Colbert (2004), we calculate theoretical offsets using
velocities estimated for a three-body interaction kick. We estimate the X-ray
binary lifetime using calculations from the previous section. As before, we
compare the theoretical offsets with the measured offsets for our ULXs, for
the two different constraints (donor age and ULX bolometric luminosity). The
results are presented in Fig. 12. For this calculation the accretion rate is
assumed to be sub-Eddington. The cluster age estimated using Starburst99
modeling and OM observations constrains the age of the binary companion, but
for the interloper we assume a fixed mass of 20 M☉. A young O-type star is
needed to create kicks large enough to eject the IMBH and, further, kick the
IMBH to the offsets observed. In addition the selection of a young interloper
with high mass is in line with our goal to estimate upper limits for the
offsets. We note that the interloper does not have to be a massive star, it
could also be a 20 M☉ BH.

We consider again Case A and Case B mass transfer. Here too, Case A assumes
that the mass transfer starts early when the donor star is on the main
sequence and Case B assumes it begins after the donor has evolved off the main
sequence. For each mass transfer case, we also look at two accretion rates:
one at the Eddington limit and one at 10% of the Eddington limit. As we did
before for natal kicks, we present the results used for the donor star
calculation and bolometric luminosity calculation in Fig. 12.

1\. Using the age estimates for the clusters (left panel of Fig. 12): For both
case A and case B mass transfer significant offsets ($>$10 pc) are possible
even for BH masses of a few 1000 M☉, and most of our data are consistent with
this scenario. On the other hand, the velocities predicted for such BHs are
smaller than 10 km s-1 (they are shown on the green line at the bottom of the
plot), implying that the system can not realistically escape the cluster. For
NGC 247 XMM1 the offset is consistent with a BH mass $<$700 M☉ for case A and
$<$2000 M☉ for case B (if it somehow escapes the cluster).

2\. Using the X-ray-based bolometric luminosities from WMR (right panel of
Figure 12): As with the stellar-mass BH scenario, only case A is considered
for the bolometric luminosity offset calculation. Just like Figure 11, the
constraint based on the luminosity suggests lower offsets than the one based
on the donor mass. Very few of the ULXs located on the left side of the plot
have offsets consistent with the bolometric luminosity predictions. Those ULXs
that do show some consistency, only correspond to models of accretion close to
the Eddington limit (dashed lines). Moreover, the BH masses for these ULXs do
not exceed 80 M☉. Again we mark the location of a kick velocity of 10 km s-1
in the plot by thick black lines. For accretion at 10% of the Eddington limit
(continuous lines) the offsets and ejection velocities are too small to
explain the data. The offset and luminosity for NGC 247 XMM1 are not
consistent with a ULX ejected 64 pc from the cluster via three-body
interactions.

##  6  DISCUSSION

###  6.1 ULXs Located Inside Star-forming Regions (Clusters)

Simulations have shown that in young and dense clusters, IMBHs can form
through core collapse, the black holes having masses of 0.1%$-$0.2% of the
cluster mass (Portegies Zwart & McMillan, 2002; Gürkan et al., 2004). We
assume therefore that IMBHs can form only in very dense clusters with masses
$>$105 M☉. One such example was thought to be M82 X-1. Portegies Zwart et al.
(2004) found an X-ray position coincident with the cluster MCG 11, the cluster
having a mass of 105.5 M☉, and a half-light radius of 1.2 pc. Recently,
however, Voss et al. (2011) showed that the ULX is actually offset by 16 pc
from the super-cluster (assuming a distance of 5.2 Mpc for M82). Supporting
the dense cluster theory further, the authors also discovered a new ULX in NGC
7479, and showed that it is a true example of a IMBH located inside a super-
cluster with mass in excess of 7$\times$105 M☉ and half-light radius of
$\sim$20 pc. Two additional ULXs in the Antennae galaxy were also shown to be
near young super-clusters (Clark et al., 2011).

Once the IMBH is formed, it still needs to capture a companion star to become
a ULX. Blecha et al. (2006) studied the probability of an IMBH becoming a ULX
using simulations of young ($<$100 Myr) clusters. For dense clusters they
found that the probability of capture through binary exchange is relatively
high, but that in general the resulting binary would appear as a ULX for only
2% of the total transfer time. This implies a typical ULX lifetime shorter
than 0.1 Myr, and a low probability to observe such a ULX. Baumgardt et al.
(2006) considered higher-mass IMBHs and found higher probabilities. Their
results show that the captured companions are massive main-sequence stars,
capable of producing high X-ray luminosities over timescales of $\sim$10 Myr.

The results in Table The Local Environment of Ultraluminous X-ray Sources
Viewed by XMM Newton’s Optical Monitor suggest that 18 of the ULXs in our
sample are located inside OB associations or SFRs. There are 3 ULXs which
apparently reside inside young SFRs according to the results in Table 4 ($<$10
Myr, Columns 6$-$9). However the cluster masses are not large enough to
contain young super-clusters in which IMBHs might grow. The estimates in Table
4 show that when detection and age estimates were possible, the majority of
the stellar populations have ages $\lesssim$20 Myr. This is true both for the
100 pc regions (col. 6) and the SFRs that contain ULXs (ULXs in SFRs are in
bold in col. 10, corresponding ages are shown in col. 11). In such young
regions, the donors are likely massive stars. In addition, super-Eddington
luminosities can theoretically be achieved with stellar-mass BHs rather than
IMBHs, as suggested by Podsiadlowski et al. (2003).

IMBHs can also form in older globular clusters through the much slower
mechanism of BH mergers (Miller & Hamilton, 2002). If a larger BH ($\gtrsim$50
M☉) is initially formed in the cluster, it can grow through mergers with
smaller black holes. Hardening of binaries through three-body interactions
cause the mergers. The process is thought to be slow and may be inefficient,
as many interlopers are ejected. O’Leary et al. (2006) built onto this
mechanism gravitational radiation and successive mergers of binary and single
BHs (Gültekin et al., 2004), as well as binary-binary interactions. In this
case, the host clusters for ULXs should be older clusters ($>$100 Myr), and
the IMBH companions later type stars. Only NGC 253 XMM1 seems to be located
inside an older stellar population according to our estimates in Table 4.
Given the resolution of the OM instrument, we cannot tell if compact clusters
are present. Moreover, the UV is not very well suited to detect old stellar
populations.

###  6.2 Nearby Clusters and the Runaway Scenario

If our ULXs are not apparently located inside super-clusters, they could still
have formed in such clusters and later been ejected. As mentioned before, it
is difficult with OM to identify dense super-clusters inside star-forming
regions. However, some of the star-forming regions located nearby (Table 4)
are young and massive enough to host such clusters. If ULXs are super-
Eddington sources, they require massive stars as companions ($\gtrsim$ 10 M☉,
type B2 or earlier, see Podsiadlowski et al., 2003). We have found that the
majority of ULXs (17, see cols. 14-17 of Table 4) are located nearby SFRs
young enough to host such massive stars. Rappaport et al. (2005) estimate that
in star-forming regions, ULXs can have a typical life of 10 to 30 Myr after a
starburst, but this lifetime can be as long as 100 Myr.

The runaway scenario is normally used as evidence for the stellar-mass BH in
ULXs, the argument being that more massive BHs are not easily ejected from
clusters. However, three-body interactions can eject even 1000 M☉ BHs. In
Section 5 we have therefore tested runaway scenarios for both stellar-mass BHs
and IMBHs. Specifically the two scenarios are: 1) stellar-mass black hole
binaries emitting at ten times the Eddington limit ejected by natal kicks, or
2) ULXs with IMBHs accreting within the Eddington limit, ejected by three-body
interactions in dense environments. We have found that the first scenario fits
the data much better (Fig. 11). However that does not completely rule out some
ULXs in this sample as examples of the latter. The left panel shows that the
IMBH scenario only works for BH masses up to $\sim$700 M☉, while in the right
panel the offsets cannot be explained by this scenario (assuming case A only).
Mapelli et al. (2011) showed with simulations that massive BHs (up to 80 M☉)
can be easily ejected via three-body encounters inside dense clusters, however
IMBHs have a low probability of being ejected because this requires a high-
mass interloper.

Finally, another scenario in which ULXs containing IMBHs are associated with
star-forming regions but not with dense clusters has been proposed by Miller &
Hamilton (2002). An IMBH formed in the core of a globular cluster might be
released in the galactic disk when the cluster disperses. If the IMBH
encounters a molecular cloud, the interaction could precipitate the collapse
of the molecular cloud and produce star formation. In the newly formed star
cluster the massive stars would sink towards the center and be captured by the
IMBH. This scenario could explain the fact that many ULXs are close to star-
forming regions but not to dense stellar clusters.

####  6.2.1 Problems with natal kicks

There are a few problems with the natal kicks scenario. Large kick velocities,
as used in this paper, can easily disrupt the binary. Even if it survives, the
orbit will become highly eccentric. Rappaport et al. (2005) and Mapelli et al.
(2011) note that such systems will have an unstable mass transfer if the mass
of the donor is larger than the BH. Moreover, the mass-transfer phase is very
short, making observing these objects in the ULX phase unlikely.

Our assumption that the accretion is super-Eddington is not very probable if
the donor is a main sequence star with a low mass. Podsiadlowski et al. (2003)
have shown that even if we assume that the accretion is super-Eddington,
luminosities $>$ 1039 erg s-1 are only obtained for donors with masses
$\gtrsim$ 10 M☉, which supports this assumption.

The natal kick velocities for BHs depend on the amount of fallback associated
with the supernova explosion. For lower-mass BHs the velocities will be
lowered with the amount of fallback (Belczynski et al., 2004). Moreover, kicks
are not expected for higher-mass BHs that form without an explosion.
Belczynski et al. (2004) estimated that the limit on the progenitor mass is
$\sim$40 M☉.

####  6.2.2 Newtonian and recoil kicks

Newtonian kicks can eject lower-mass IMBHs from clusters that are not very
dense (with escape velocities $\lesssim$ 20 km s-1, as discussed in Section
5.2). However IMBHs can be subject to higher velocity, recoil kicks during BH
mergers. Recently, important advances in numerical relativity have allowed
precise calculations of the dynamics of BH mergers. These results show that
the asymmetry during the plunge phase of the inspiral can produce “recoil
kicks” with velocities as high as 4000 km s-1 (e.g. Baker et al., 2008). The
kick velocities depend on the mass ratio and spins of the black holes. The
velocities are largest when the BHs are of comparable mass and both have large
spins. These large kicks have very important implications on IMBH growth in
clusters. Holley-Bockelmann et al. (2008) showed that IMBHs with masses $<$
500 M☉ are very easily ejected, unless the merging companion BH has a mass $<$
10M☉. Baker et al. (2008) obtained different results, however. They found that
even a $\sim$ 200 M☉ BH can sustain radiative recoil kicks from a merger and
never leave the cluster. This is the same limiting mass that Newtonian three-
body kicks can eject (O’Leary et al., 2006). Jonker et al. (2010) was the
first to propose this scenario for the ULX, CXO J122418.6$+$144545.

A slightly different scenario for some ULXs is a recoiling supermassive BH.
Fujita (2009) showed that accretion from the interstellar medium in the
galactic disk can easily generate luminosities similar to ULXs.

##  7 CONCLUSIONS

The summary of the main results presented in this paper is as follows:

  * •

Using UVW1 photometry, we have found that 7 sources with UVW1 magnitudes
measured from 100 pc regions have emission that is consistent with a single
star (see also Table The Local Environment of Ultraluminous X-ray Sources
Viewed by XMM Newton’s Optical Monitor, col. 11 and Figure 3). Indeed for four
of these objects our UVW1-predicted spectral types are identical to those
found in the literature using other methods. We note that all of these objects
(even those with consistent published data in the literature) may be
contaminated by or dominated by an accretion disk component. Indeed for at
least two sources the UV emission likely comes predominantly from an accretion
disk (see also bullet 2).

  * •

By comparing the theoretical UV fluxes expected from accretion disks with the
UVW1 measurements, we found that, for most of the ULXs in our sample (esp. NGC
1313 XMM3 and Holmberg IX XMM1), a significant part of the measured UV flux
could come from accretion disk emission (Fig. 3). However, even if the
accretion disk is at its brightest (e.g. irradiated), it is not bright enough
in most cases to be responsible for all of the UV emission. In those cases
(the cases in Table 2, column 11 where the flux is consistent with multiple
sources), it is likely that a star forming region, an accretion disk, and a
donor star contaminates the aperture, none of which can be isolated using the
UVW1 photometry alone.

  * •

We looked at 3-5 SFRs around most ULXs. We derived statistics for the closest
SFRs to the ULX, the youngest SFRs around the ULX, and the most massive SFRs
around the ULX. Roughly half of the closest SFRs actually overlap with the
ULX. Of the youngest star forming regions, 17 out of 21 are less than 10 Myr
in age, which might imply that young regions are intrinsically related to
ULXs. Finally there are also 17 sources with at least one region that is more
massive than 105 M⊙.

  * •

OM color-color plots show relatively little reddening for most of the SFRs
close to our ULXs, implying that they are likely not IMBHs accreting from
molecular clouds. The extinction is also, in general, much less than suggested
by the absorption columns obtained from X-ray modeling, indicating that the
X-ray absorption is located close to the ULX.

  * •

We tested runaway scenarios for both stellar-mass BHs and IMBHs. Specifically
the two scenarios are: 1) stellar-mass black hole binaries emitting at ten
times the Eddington limit ejected by natal kicks, or 2) ULXs with IMBHs
accreting within the Eddington limit, ejected by three-body interactions in
dense environments. We have found that the first scenario fits the data best
(Fig. 11). On the other hand, that does not completely rule out that some ULXs
in this sample are IMBHs, though the masses we found in this analysis were
somewhat small.

We would like to thank R. Mushotzky and K. Kuntz for helpful discussions and
suggestions. We are also appreciate the helpful comments from the anonymous
referee, which have significantly improved the quality and clarity of this
paper. J. T. gratefully acknowledges support from the Hodson Trust Fellowship
Program at St. John’s College. This research has made use of the NASA/IPAC
Extragalactic Database (NED) which is operated by the Jet Propulsion
Laboratory, California Institute of Technology, under contract with the
National Aeronautics and Space Administration.

## References

  * Arzoumanian et al. (2002) Arzoumanian, Z., Chernoff, D. F., & Cordes, J. M. 2002, ApJ, 568, 289
  * Baumgardt et al. (2006) Baumgardt, H., Hopman, C., Portegies Zwart, S., & Makino, J. 2006, MNRAS, 372, 467
  * Begelman (2002) Begelman, M. C. 2002, ApJ, 568, L97
  * Belczynski et al. (2004) Belczynski, K., Sadowski, A., & Rasio, F. A. 2004, ApJ, 611, 1068
  * Berghea et al. (2008) Berghea, C. T., Weaver, K. A., Colbert, E. J. M., & Roberts, T. P. 2008, ApJ, 687, 471
  * Berghea et al. (2010a) Berghea, C. T., Dudik, R. P., Weaver, K. A., & Kallman, T. R. 2010a, ApJ, 708, 354
  * Berghea et al. (2010b) Berghea, C. T., Dudik, R. P., Weaver, K. A., & Kallman, T. R. 2010b, ApJ, 708, 364
  * Berghea & Dudik (2012) Berghea, C. T., & Dudik, R. P. 2012, ApJ, 751, 104
  * Baker et al. (2008) Baker, J. G., Boggs, W. D., Centrella, J., et al. 2008, ApJ, 682, L29
  * Blecha et al. (2006) Blecha, L., Ivanova, N., Kalogera, V., Belczynski, K., Fregeau, J., & Rasio, F. 2006, ApJ, 642, 427
  * Bodaghee et al. (2012) Bodaghee, A., Tomsick, J. A., Rodriguez, J., & James, J. B. 2012, ApJ, 744, 108
  * Böker, van der Marel, & Vacca (1999) Böker T., van der Marel R., Vacca W. D., 1999, AJ, 118, 831
  * Brassington et al. (2005) Brassington, N. J., Read, A. M., & Ponman, T. J. 2005, MNRAS, 360, 801
  * Cardelli, Clayton, & Mathis (1989) Cardelli, J.A., Clayton, G.C., & Mathis, J.S. 1989, ApJ, 345, 245
  * Clark et al. (2007) Clark, D. M., et al. 2007, ApJ, 658, 319
  * Clark et al. (2011) Clark, D. M., Eikenberry, S. S., Brandl, B. R., et al. 2011, MNRAS, 410, 890
  * Colbert et al. (1995) Colbert, E. J. M., Petre, R., Schlegel, E. M., & Ryder, S. D. 1995, ApJ, 446, 177
  * Colbert & Mushotzky (1999) Colbert, E. J. M., & Mushotzky, R. F. 1999, ApJ, 519, 89
  * Colbert & Ptak (2002) Colbert, E. J. M., & Ptak, A. F. 2002, ApJS, 143, 25
  * Coleiro & Chaty (2013) Coleiro, A., & Chaty, S. 2013, ApJ, 764, 185
  * Copperwheat et al. (2005) Copperwheat, C., Cropper, M., Soria, R., & Wu, K. 2005, MNRAS, 362, 79
  * Cordes & Chernoff (1998) Cordes, J. M., & Chernoff, D. F. 1998, ApJ, 505, 315
  * Dewangan, et al. (2005) Dewangan, G. C., Griffiths, R. E.; Choudhury, M., Miyaji, & T., Schurch, N. J. ApJ, 635, 198
  * Eracleous et al. (2002) Eracleous, M., Shields, J. C., Chartas, G., & Moran, E. C. 2002, ApJ, 565, 108
  * Feng & Kaaret (2008) Feng, H., & Kaaret, P. 2008, ApJ, 675, 1067
  * Fryer & Kalogera (2001) Fryer, C. L., & Kalogera, V. 2001, ApJ, 554, 548
  * Fujita (2009) Fujita, Y. 2009, ApJ, 691, 1050
  * Gladstone et al. (2013) Gladstone, J. C., Copperwheat, C., Heinke, C. O., et al. 2013, arXiv:1303.1213
  * Grimm et al. (2003) Grimm, H.-J., Gilfanov, M., & Sunyaev, R. 2003, MNRAS, 339, 793
  * Grisé, Pakull, & Motch (2006) Grisé, F., Pakull, M. W., & Motch, C. 2006, IAUS, 230, 302
  * Grisé et al. (2011) Grisé, F., Kaaret, P., Pakull, M. W., & Motch, C. 2011, ApJ, 734, 23
  * Grisé et al. (2012) Grisé, F., Kaaret, P., Corbel, S., et al. 2012, ApJ, 745, 123
  * Gültekin et al. (2004) Gültekin, K., Miller, M. C., & Hamilton, D. P. 2004, ApJ, 616, 221
  * Gürkan et al. (2004) Gürkan, M. A., Freitag, M., & Rasio, F. A. 2004, ApJ, 604, 632
  * Gürkan et al. (2006) Gürkan, M. A., Fregeau, J. M., & Rasio, F. A. 2006, ApJ, 640, L39
  * Han et al. (2012) Han, X., An, T., Wang, J.-Y., et al. 2012, Research in Astronomy and Astrophysics, 12, 1597
  * Hobbs et al. (2005) Hobbs, G., Lorimer, D. R., Lyne, A. G., & Kramer, M. 2005, MNRAS, 360, 974
  * Holley-Bockelmann et al. (2008) Holley-Bockelmann, K., Gultekin, K., Shoemaker, D., & Yunes, N. 2008, ApJ, 686 829
  * Humphrey et al. (2003) Humphrey, P. J., Fabbiano, G., Elvis, M., Church, M. J., & Bałucińska-Church, M. 2003, MNRAS, 344, 134
  * Irwin, Bregman & Athey (2004) Irwin, J. A., Bregman, J. N., Athey, A. E. 2004, ApJ, 601, L143
  * Jenkins et al. (2004) Jenkins, L. P., Roberts, T. P., Warwick, R. S., Kilgard, R. E., & Ward, M. J. 2004, MNRAS, 349, 404
  * Jonker et al. (2010) Jonker, P. G., Torres, M. A. P., Fabian, A. C., et al. 2010, MNRAS, 407, 645
  * Kaaret et al. (2004) Kaaret, P., Alonso-Herrero, A., Gallagher, J. S., Fabbiano, G., Zezas, A., & Rieke, M. J. 2004, MNRAS, 348, L28
  * Kaaret et al. (2006) Kaaret, P., Simet, M. G., & Lang, C. C. 2006, ApJ, 646, 174
  * Kaaret & Feng (2007) Kaaret, P., & Feng, H. 2007, ApJ, 669, 106
  * Kaaret et al. (2010) Kaaret, P., Feng, H., Wong, D. S., & Tao, L. 2010, ApJ, 714, L167
  * Kong (2003) Kong, A. K. H. 2003, MNRAS, 346, 265
  * Krolik (2004) Krolik, J. H. 2004, ApJ, 615, 383
  * Kuntz et al. (2005) Kuntz, K.D. et al. 2005, ApJ, 620, L31
  * Kuntz et al. (2008) Kuntz, K. D., Harrus, I., McGlynn, T. A., Mushotzky, R. F., & Snowden, S. L. 2008, PASP, 120, 740
  * Lang et al. (2007) Lang, C. C., Kaaret, P., Corbel, S. & Mercer, A. 2007, ApJ, 666, 79
  * Lehmann et al. (2005) Lehmann, I., et al. 2005, A&A, 431, 847
  * Leitherer et al. (1999) Leitherer, C. et al. 1999, ApJS, 123, 3
  * Lira et al. (2000) Lira, P., Lawrence, A., & Johnson, R. A. 2000, MNRAS, 319, 17
  * Liu et al. (2002) Liu, J.-F., Bregman, J.N., & Seitzer, P. 2002, ApJ, 580, L31
  * Liu et al. (2004) Liu, J.-F., Bregman, J.N., & Seitzer, P. 2004, ApJ, 602, 249
  * Liu et al. (2007) Liu, J.-F., Bregman, J., Miller, J., & Kaaret, P. 2007, ApJ, 661, 165
  * Liu et al. (2009) Liu, J., Bregman, J. N., & McClintock, J. E. 2009, ApJ, 690, L39
  * Liu (2009) Liu, J. 2009, ApJ, 704, 1628
  * Mapelli et al. (2009) Mapelli, M., Colpi, M., & Zampieri, L. 2009, MNRAS, 395, L71
  * Mapelli et al. (2011) Mapelli, M., Ripamonti, E., Zampieri, L., & Colpi, M. 2011, MNRAS, 416, 1756
  * Miller & Hamilton (2002) Miller, M. C., & Hamilton, D. P. 2002, MNRAS, 330, 232
  * Miller et al. (2004) Miller, J. M., Fabian, A. C., & Miller, M. C. 2004, ApJ, 607, 931
  * Miller & Colbert (2004) Miller, M. C., & Colbert, E. J. M. 2004, International Journal of Modern Physics D, 13, 1
  * Miller-Jones et al. (2009) Miller-Jones, J. C. A., Jonker, P. G., Dhawan, V., et al. 2009, ApJ, 706, L230
  * Mineo et al. (2013) Mineo, S., Rappaport, S., Steinhorn, B., et al. 2013, arXiv:1301.4084
  * Moon et al. (2011) Moon, D.-S., Harrison, F. A., Cenko, S. B., & Shariff, J. A. 2011, ApJ, 731, L32
  * Motch et al. (2011) Motch, C., Pakull, M. W., Grisé, F., & Soria, R. 2011, Astronomische Nachrichten, 332, 367
  * O’Leary et al. (2006) O’Leary, R. M., Rasio, F. A., Fregeau, J. M., Ivanova, N., & O’Shaughnessy, R. 2006, ApJ, 637, 937
  * Pakull et al. (2010) Pakull, M. W., Soria, R., & Motch, C. 2010, Nature, 466, 209
  * Patruno & Zampieri (2008) Patruno, A., & Zampieri, L. 2008, MNRAS, 386, 543
  * Pintore & Zampieri (2012) Pintore, F., & Zampieri, L. 2012, MNRAS, 420, 1107
  * Podsiadlowski et al. (2003) Podsiadlowski, P., Rappaport, S., & Han, Z. 2003, MNRAS, 341, 385
  * Portegies Zwart et al. (1999) Portegies Zwart, S. F., Makino, J., McMillan, S. L. W., & Hut, P. 1999, A&A, 348, 117
  * Portegies Zwart & McMillan (2002) Portegies Zwart, S. F., & McMillan, S. L. W. 2002, ApJ, 576, 899
  * Portegies Zwart et al. (2004) Portegies Zwart, S. F., Baumgardt, H., Hut, P., Makino, J., & McMillan, S. L. W. 2004, Nature, 428, 724
  * Poutanen et al. (2012) Poutanen, J., Fabrika, S., Valeev, A. F., Sholukhova, O., & Greiner, J. 2012, arXiv:1210.1210
  * Prestwich et al. (2012) Prestwich, A. H., Galache, J. L., Linden, T., et al. 2012, ApJ, 747, 150
  * Ptak et al. (2006) Ptak, A., Colbert, E., van der Marel, R. P., Roye, E., Heckman, T., & Towne, B. 2006, ApJS, 166, 154
  * Ramsey et al. (2006) Ramsey, C. J., Williams, R. M., Gruendl, R. A., et al. 2006, ApJ, 641, 241
  * Rangelov et al. (2011) Rangelov, B., Prestwich, A. H., & Chandar, R. 2011, ApJ, 741, 86
  * Rangelov et al. (2012) Rangelov, B., Chandar, R., Prestwich, A., & Whitmore, B. C. 2012, ApJ, 758, 99
  * Ranalli, Comastri, & Setti (2003) Ranalli, P., Comastri, A., & Setti, G. 2003, A&A, 399, 39
  * Rappaport et al. (2005) Rappaport, S. A., Podsiadlowski, P., & Pfahl, E. 2005, MNRAS, 356, 401
  * Rappaport et al. (2010) Rappaport, S., Levine, A., Pooley, D., & Steinhorn, B. 2010, ApJ, 721, 1348
  * Repetto et al. (2012) Repetto, S., Davies, M. B., & Sigurdsson, S. 2012, MNRAS, 425, 2799
  * Reid et al. (2011) Reid, M. J., McClintock, J. E., Narayan, R., et al. 2011, ApJ, 742, 83
  * Ripamonti et al. (2011) Ripamonti, E., Mapelli, M., Zampieri, L., & Colpi, M. 2011, Astronomische Nachrichten, 332, 418
  * Roberts & Warwick (2000) Roberts, T. P., & Warwick, R. S. 2000, MNRAS, 315, 98
  * Roberts et al. (2002) Roberts, T. P., Warwick, R. S., Ward, M. J., & Murray, S. S. 2002, MNRAS, 337, 677
  * Roberts et al. (2008) Roberts, T. P., Levan, A. J., & Goad, M. R. 2008, MNRAS, 387, 73
  * Roberts et al. (2011) Roberts, T. P., Gladstone, J. C., Goulding, A. D., et al. 2011, Astronomische Nachrichten, 332, 398
  * Roberts et al. (2012) Roberts, T. P., Fabbiano, G., Luo, B., et al. 2012, ApJ, 760, 135
  * Sanbuichi et al. (1993) Sanbuichi, K., Yamada, T. T., & Fukue, J. 1993, PASJ, 45, 443
  * Sepinsky et al. (2005) Sepinsky, J., Kalogera, V., & Belczynski, K. 2005, ApJ, 621, L37
  * Shakura & Sunyaev (1973) Shakura, N. I., & Sunyaev, R. A. 1973, A&A, 24, 337
  * Soria et al. (2005) Soria, R., Cropper, M., Pakull, M., Mushotzky, R., & Wu, K. 2005, MNRAS, 356, 12
  * Soria et al. (2006) Soria, R., Fender, R. P., Hannikainen, D. C., Read, A. M., & Stevens, I. R. 2006, MNRAS, 368, 1527
  * Soria et al. (2012) Soria, R., Hakala, P. J., Hau, G. K. T., Gladstone, J. C., & Kong, A. K. H. 2012, MNRAS, 420, 3599
  * Swartz et al. (2003) Swartz, D. A., Ghosh, K. K., McCollough, M. L., Pannuti, T. G., Tennant, A. F.; Wu, K. 2003, ApJS, 144, 213
  * Swartz et al. (2009) Swartz, D. A., Tennant, A. F., & Soria, R. 2009, ApJ, 703, 159
  * Strohmayer (2009) Strohmayer, T. E. 2009, ApJ, 706, L210
  * Tao et al. (2011) Tao, L., Feng, H., Grisé, F., & Kaaret, P. 2011, ApJ, 737, 81
  * Tao et al. (2012a) Tao, L., Feng, H., Kaaret, P., Grisé, F., & Jin, J. 2012a, ApJ, 758, 85
  * Tao et al. (2012b) Tao, L., Kaaret, P., Feng, H., & Grisé, F. 2012b, ApJ, 750, 110
  * Terashima et al. (2006) Terashima, Y., Inoue, H., & Wilson, A. 2006, ApJ, 645, 264
  * van Paradijs & McClintock (1994) van Paradijs, J., & McClintock, J. E. 1994, A&A, 290, 133
  * Vogler & Pietsch (1999) Vogler, A., & Pietsch, W. 1999, A&A, 342, 101
  * Voss et al. (2011) Voss, R., Nielsen, M. T. B., Nelemans, G., Fraser, M., & Smartt, S. J. 2011, MNRAS, 418, L124
  * Winter et al. (2006) Winter, L.M., Mushotzky, R.F., & Reynolds, C.S. 2006, ApJ, 649, 730 (WMR)
  * Winter et al. (2007) Winter, L.M., Mushotzky, R.F., & Reynolds, C.S. 2007, ApJ, 655, 163
  * Yang et al. (2011) Yang, L., Feng, H., & Kaaret, P. 2011, ApJ, 733, 118
  * Zampieri et al. (2004) Zampieri, L., Mucciarelli, P., Falomo, R., Kaaret, P., Di Stefano, R., Turolla, R., Chieregato, M., & Treves, A. 2004, ApJ, 603, 523
  * Zampieri et al. (2012) Zampieri, L., Impiombato, D., Falomo, R., Grisé, F., & Soria, R. 2012, MNRAS, 419, 1331
  * Zezas et al. (2002) Zezas, A., Fabbiano, G., Rots., A. H., & Murray, S. S. 2002, ApJ, 577, 710
  * Zezas & Fabbiano (2002) Zezas, A., & Fabbiano, G. 2002, ApJ, 577, 726
  * Zuo & Li (2010) Zuo, Z.-Y., & Li, X.-D. 2010, MNRAS, 405, 2768

Table 1: OM observations

Gal | D | AV | Obs ID | Date | Filters |
---|---|---|---|---|---|---
| (Mpc) |  |  | (s) |  |
(1) | (2) | (3) | (4) | (5) | (6) |
NGC 247 | 3.09 | 0.06 | 0601010101 | 2009 Dec 27 | U, B, UVW1, UVM2 |
NGC 253 | 3.73 | 0.06 | 0304850901 | 2006 Jan 2 | U |
$\cdots$ | $\cdots$ | $\cdots$ | 0110900101 | 2000 Dec 13 | B, UVW1, UVM2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0152020101 | 2003 Jun 19 | B, UVW1, UVW2 |
NGC 1313 | 4.17 | 0.34 | 0106860101 | 2000 Oct 17 | V, UVW1 |
$\cdots$ | $\cdots$ | $\cdots$ | 0405090101 | 2006 Oct 15 | B, UVW1, UVW2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0150280701 | 2003 Dec 27 | U, UVM2 |
IC 0342 | 3.90 | 1.73 | 0093640901 | 2001 Feb 11 | UVW1 |
$\cdots$ | $\cdots$ | $\cdots$ | 0206890101 | 2004 Feb 20 | B, UVM2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0206890201 | 2004 Aug 17 | U, UVW1, UVM2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0206890401 | 2005 Feb 10 | V, UVW1, UVM2 |
NGC 2403 | 3.56 | 0.12 | 0150651101 | 2003 Apr 30 | UVW2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0150651201 | 2003 Sep 11 | U |
$\cdots$ | $\cdots$ | $\cdots$ | 0164560901 | 2004 Sep 12 | UVW1, UVM2 |
Holmberg II | 2.70 | 0.10 | 0112520701 | 2002 Apr 16 | UVW1 |
$\cdots$ | $\cdots$ | $\cdots$ | 0561580401 | 2010 Mar 26 | UVW1 |
Holmberg IX | 3.60 | 0.24 | 0200980101 | 2004 Sep 26 | UVW1, UVM2, UVW2 |
NGC 4395 | 4.00 | 0.05 | 0142830101 | 2003 Nov 30 | U, B, UVW1, UVW2 |
NGC 4490 | 7.80 | 0.07 | 0112280201 | 2002 May 27 | V |
$\cdots$ | $\cdots$ | $\cdots$ | 0556300201 | 2008 Jun 22 | U, UVW1, UVM2 |
NGC 4631 | 7.50 | 0.05 | 0110900201 | 2002 Jun 28 | U, B, UVW1, UVW2 |
NGC 4736 | 4.30 | 0.05 | 0094360601 | 2002 May 23 | V, B |
$\cdots$ | $\cdots$ | $\cdots$ | 0404980101 | 2006 Nov 27 | U |
NGC 5204 | 4.80 | 0.04 | 0142770101 | 2003 Jan 6 | UVM2, UVW2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0150650301 | 2003 May 1 | B |
$\cdots$ | $\cdots$ | $\cdots$ | 0405690201 | 2006 Nov 19 | UVW1 |
M51 | 7.20 | 0.11 | 0303420101 | 2006 May 20 | U, UVW1 |
$\cdots$ | $\cdots$ | $\cdots$ | 0303420201 | 2006 May 24 | UVW2 |
M101 | 7.40 | 0.03 | 0104260101 | 2002 Jun 4 | UVW1, UVM2, UVW2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0212480201 | 2005 Jan 8 | V, U, B |
$\cdots$ | $\cdots$ | $\cdots$ | 0164560701 | 2004 Jul 23 | U, UVW1, UVM2, UVW2 |
NGC 5408 | 4.80 | 0.21 | 0500750101 | 2008 Jan 13 | UVW2 |
$\cdots$ | $\cdots$ | $\cdots$ | 0653380301 | 2010 Jul 19 | UVM2 |
NGC 6946 | 5.10 | 1.06 | 0093641701 | 2003 Jun 18 | UVW1 |
$\cdots$ | $\cdots$ | $\cdots$ | 0200670101 | 2004 Jun 9 | U |
$\cdots$ | $\cdots$ | $\cdots$ | 0200670401 | 2004 Jun 25 | B, UVM2 |

Note. — (1) Host galaxy; (2) Galaxy distance adopted from WMR; (3) Galactic
extinction in the direction of each galaxy, see text for details; (4) XMM-
Newton Observation ID; (5) Observation date; (6) OM filters used in this paper
for each observation;

Table 2: ULX sample and photometric results for the 100 pc regions

Src | IAU | ULX | RA | DEC | Exp | LX | UVW1 | S/N | FX/ | Stars | Literature | Names | Ref
---|---|---|---|---|---|---|---|---|---|---|---|---|---
| (J2000) |  | (J2000) | (J2000) | (Ks) | (erg/s) |  |  | FUVW1 |  |  |  |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10) | (11) | (12) | (13) | (14)
1 | 0047-2047 | N247 XMM1 | 00 47 03.8 | -20 47 46.2 | 4.4 | 39.85 | 17.59 $\pm$ 0.04 | 22.6 | 2.10 | 5.2 | B1-B7 Ib1 | X1 | 27
2 | 0047-2517 | N253 XMM1 | 00 47 32.8 | -25 17 52.6 | 3.2 | 39.71 | 15.68 $\pm$ 0.04 | 20.3 | 1.71 | 43.7 | $\cdots$ | PSX-2, X33, C5 | 11, 12, 13
3 | 0047-2520 | N253 XMM2 | 00 47 22.4 | -25 20 55.2 | 3.2 | 39.43 | 17.69 $\pm$ 0.07 | 10.7 | 1.68 | 6.8 | $\cdots$ | PSX-5, X21, C4 | 11, 12, 13
4 | 0047-2015 | N253 XMM6 | 00 47 42.8 | -25 15 05.5 | 2.0 | 39.49 |  $>$ 19.09 - | $\cdots$ |  $>$ 3.06 |  $<$ O8V | $\cdots$ | X40 | 12
5 | 0318-6636 | N1313 XMM3 | 03 18 22.5 | -66 36 06.2 | 4.0 | 40.34 | 20.41 $\pm$ 0.15 | 5.3 | 3.07 | O9V | B2-O72 | X-2, U5 | 14, 13
6 | 0345+6804 | IC342 XMM1 | 03 45 55.8 | +68 04 54.5 | 4.0 | 39.80 | 18.19 $\pm$ 0.12 | 4.5 | 2.76 | 4.7 | F8-G0 Ib3 | X-1, U6 | 15, 13
7 | 0346+6811 | IC342 XMM2 | 03 46 15.0 | +68 11 11.2 | 4.0 | 39.93 |  $>$ 18.63 - | $\cdots$ |  $>$ 2.10 |  $<$ 3.1 | $\cdots$ | X-13 | 15
8 | 0346+6805 | IC342 XMM3 | 03 46 48.6 | +68 05 43.2 | 4.0 | 40.75 | 8.40 $\pm$ 0.00 | 1461.0 | 0.56 | 39026.0 | $\cdots$ | X-21 | 15
9 | 0736+6535 | N2403 XMM1 | 07 36 25.6 | +65 35 40.0 | 4.3 | 39.49 | 19.10 $\pm$ 0.05 | 19.3 | 1.74 | O3V | OV/BI4 | X-1, U7 | 16, 13
10 | 0819+7042 | Ho II XMM1 | 08 19 28.8 | +70 42 20.3 | 4.4 | 40.00 | 16.73 $\pm$ 0.02 | 134.9 | 1.52 | 8.6 | B2Ib-O4V5 | X-1, IXO 31 | 17, 18
11 | 0957+6903 | Ho IX XMM1 | 09 57 53.3 | +69 03 48.7 | 7.9 | 40.20 | 20.39 $\pm$ 0.06 | 16.1 | 2.92 | B2III |  $>$ B2 (I/III)6 | Ho X-1, M81 X-9 | 19
12 | 1226+3331 | N4395 XMM1 | 12 26 01.5 | +33 31 29.0 | 4.3 | 39.43 | 19.33 $\pm$ 0.06 | 22.5 | 1.62 | O3V | $\cdots$ | X-1, IXO 53 | 16, 18
13 | 1230+4139 | N4490 XMM1 | 12 30 32.4 | +41 39 14.6 | 4.3 | 39.81 | 19.17 $\pm$ 0.04 | 45.7 | 1.32 | 7.6 | $\cdots$ | ULX4 | 20
14 | 1230+4138 | N4490 XMM2 | 12 30 36.5 | +41 38 33.3 | 4.3 | 39.67 | 17.84 $\pm$ 0.02 | 63.7 | 0.64 | 26.0 | $\cdots$ | ULX5 | 20
15 | 1230+4138 | N4490 XMM3 | 12 30 43.3 | +41 38 11.5 | 4.3 | 40.94 | 19.02 $\pm$ 0.04 | 55.1 | 2.39 | 8.8 | $\cdots$ | ULX6 | 20
16 | 1230+4139 | N4490 XMM4 | 12 30 31.1 | +41 39 08.1 | 4.3 | 39.79 | 17.00 $\pm$ 0.03 | 43.1 | 1.03 | 56.3 | $\cdots$ | ULX3 | 20
17 | 1230+4141 | N4490 XMM5 | 12 30 30.3 | +41 41 40.3 | 4.3 | 39.47 | 17.78 $\pm$ 0.04 | 54.2 | 0.94 | 27.7 | $\cdots$ | ULX2 | 20
18 | 1241+3232 | N4631 XMM1 | 12 41 55.8 | +32 32 14.0 | 2.3 | 39.81 | 16.83 $\pm$ 0.03 | 77.9 | 0.73 | 61.4 | $\cdots$ | X-1, U29, IXO 65 | 16, 13, 17
19 | 1250+4107 | N4736 XMM1 | 12 50 50.2 | +41 07 12.0 | 4.4 | 40.25 | 13.69 $\pm$ 0.01 | 175.1 | 1.17 | 747.3 | $\cdots$ | X-4 | 21
20 | 1329+5825 | N5204 XMM1 | 13 29 38.5 | +58 25 03.6 | 5.0 | 39.83 | 18.07 $\pm$ 0.03 | 76.5 | 1.38 | 8.0 | B0Ib7 | X-1, U36, IXO 77 | 16, 13, 17
21 | 1330+4711 | M51 XMM2 | 13 30 07.7 | +47 11 04.8 | 4.0 | 39.48 | 19.12 $\pm$ 0.08 | 13.4 | 1.55 | 6.9 | F2-5I8 | ULX 9, IXO 81, 82 | 22, 18, 23
22 | 1329+4710 | M51 XMM6 | 13 29 57.5 | +47 10 45.3 | 4.0 | 40.54 | 17.34 $\pm$ 0.04 | 53.4 | 1.89 | 35.3 | $\cdots$ | 63 | 23
23 | 1403+5418 | M101 XMM1 | 14 03 14.7 | +54 18 05.0 | 2.0 | 39.46 | 20.72 $\pm$ 0.17 | 5.1 | 1.79 | O6V | $\cdots$ | U39, XMM-2 | 13, 24
24 | 1403+5427 | M101 XMM2 | 14 03 03.8 | +54 27 37.0 | 1.3 | 39.66 | 20.82 $\pm$ 0.20 | 4.3 | 2.10 | O7V | $\cdots$ | U38, XMM-1 | 13, 24
26 | 1404+5426 | M101 XMM3 | 14 04 14.6 | +54 26 04.4 | 1.3 | 39.57 | 21.04 $\pm$ 0.22 | 3.8 | 2.08 | B5I | B0I9 | U41, XMM-3 | 13, 24
25 | 1403-4122 | N5408 XMM1 | 14 03 19.8 | -41 22 59.3 | 4.0 | 40.04 | 17.39 $\pm$ 0.05 | 37.2 | 1.50 | 8.3 | $\cdots$ | X-1 | 25
27 | 2035+6011 | N6946 XMM1 | 20 35 00.8 | +60 11 30.6 | 1.9 | 40.00 | 19.14 $\pm$ 0.14 | 5.1 | 2.56 | 3.4 | O/BI10 | X-1, U45, IXO 85 | 26, 13, 17

Note. — NGC 4736 XMM1 and NGC 5408 XMM1 have no images in the UVW1 band. The
values shown for these ULXs were obtained using the U and UVM2 band. (2) IAU
names; (3) ULX names from WMR; (6) Exposure time for the UVW1 image used; (7)
The logarithm of the X-ray unabsorbed luminosities for ULXs in the 0.3$-$10
keV band from WMR; (8) UVW1 magnitude estimated by integrating the fluxes
within the 100 pc regions. For NGC 253 XMM6 and IC 342 XMM2 we show the
3$\sigma$ lower limit; (9) Signal to noise ratio in the UVW1 band; (10) The
logarithm of the ratio between the X-ray and UVW1 luminosities; (11) This
column shows estimates for the stars detected within the 100 pc regions.
Assuming the UVW1 emission is due to young stars we estimate the minimum
number of O5 V stars required to produce an equivalent flux in the UVW1 band.
For the estimates compatible with single stars we show the star types based
only on the UVW1 flux. Four of these we estimate to have actually single star
counterparts (Section 3.1), and their types are printed in bold. (12) The star
types for known optical counterparts from the literature; (13) Common names
for ULXs from the literature; (14) References for the names in col. 13.

References. — 1\. Tao et al. (2012a); 2\. Liu et al. (2007); 3\. Feng & Kaaret
(2008); 4\. Roberts et al. (2008); 5\. Kaaret et al. (2004); 6\. Grisé,
Pakull, & Motch (2006); 7\. Liu et al. (2004); 8\. Terashima et al. (2006);
9\. Grisé et al. (2012); 10\. Kaaret et al. (2010); Berghea & Dudik (2012);
11\. Humphrey et al. (2003); 12\. Vogler & Pietsch (1999); 13\. Berghea et al.
(2008); 14\. Colbert et al. (1995); 15\. Kong (2003); 16\. Roberts & Warwick
(2000); 17\. Lehmann et al. (2005); 18\. Colbert & Ptak (2002); 19\. Miller et
al. (2004); 20\. Roberts et al. (2002); 21\. Eracleous et al. (2002); 22\.
Terashima et al. (2006); 23\. Dewangan, et al. (2005); 24\. Jenkins et al.
(2004); 25\. Soria et al. (2006); 26\. Berghea & Dudik (2012). 27\. Lira et
al. (2000);

Table 3: ULXs with multiple observations

ULX | Date | UT | Mag
---|---|---|---
|  | (h) |
(1) | (2) | (3) | (4)
Holmberg II XMM1 | 2002 Apr 16 | 12.171 | 16.735 $\pm$ 0.021
$\cdots$ | 2010 Mar 26 | 9.499 | 16.825 $\pm$ 0.029
$\cdots$ | 2010 Mar 26 | 11.612 | 16.995 $\pm$ 0.017
$\cdots$ | 2010 Mar 26 | 17.837 | 16.875 $\pm$ 0.015
IC 342 XMM1 | 2004 Aug 17 | 21.629 | 18.195 $\pm$ 0.115
$\cdots$ | 2005 Feb 10 | 20.117 | 18.225 $\pm$ 0.117
$\cdots$ | 2001 Feb 11 | 1.345 |  $<$ 18.21
M101 XMM2 | 2004 Jul 23 | 10.751 | 21.037 $\pm$ 0.255
$\cdots$ | 2005 Jan 8 | 16.396 | 21.297 $\pm$ 0.203
NGC 1313 XMM3 | 2000 Oct 17 | 8.269 | 20.405 $\pm$ 0.153
$\cdots$ | 2006 Oct 16 | 13.782 | 20.425 $\pm$ 0.111

Note. — The data in the last column is for the UVW1 band, except for M101
XMM2, which is the U band.

Table 4: Star-forming regions estimates

ULX | Average Values | 100 pc regions | Closest SFR | Youngest SFR | Most massive SFR
---|---|---|---|---|---
ULX | AV(X) | AV | AV | Age | Age | Star | AV | Mass | Offset | Age | Star | Mass | Offset | Age | Star | Mass | Offset | Age | Star | Mass
|  | Range | Avg. | Myr | Myr |  |  | log M☉ | pc | Myr |  | log M☉ | pc | Myr |  | log M☉ | pc | Myr |  | log M☉
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10) | (11) | (12) | (13) | (14) | (15) | (16) | (17) | (18) | (19) | (20) | (21)
N247 XMM1 | 2.31 | 0$-$1.5 | 0.83 | 18 | 10 | B0.5 | 1.0 | 3.8 | 64 | 10 | B0.5 | 3.8 | 492 | 3 | O6 | 3.8 | 190 | 10 | B0.5 | 5.2
N253 XMM1 | 2.82 | 1$-$2.5 | 1.75 | 500 | 900 | A5 | 1.0 | 6.3 | 211 | 100 | B4.5 | 6.9 | 309 | 100 | B4.5 | 6.9 | 309 | 100 | B4.5 | 6.9
N253 XMM2 | 1.01 | 1.00 | 1.00 | 13 | 100 | B4.5 | 1.0 | 4.8 | 60 | 8 | B0 | 4.1 | 60 | 8 | B0 | 4.1 | 462 | 15 | B1.5 | 4.7
N253 XMM6 | 3.55 | 1$-$1.2 | 0.73 | 118 | $\cdots$ | $\cdots$ | $\cdots$ |  $<$3.4 | 221 | 5 | O9 | 4.6 | 221 | 5 | O9 | 4.6 | 221 | 150 | B5.5 | 5.1
N1313 XMM3 | 3.49 | 0.50 | 0.50 | 5 | $\cdots$ | $\cdots$ | $\cdots$ | 3.1 | 68 | $\cdots$ | $\cdots$ | 3.2 | 106 | 5 | O9 | 3.0 | 222 | $\cdots$ | $\cdots$ | 3.6
IC342 XMM1 | 3.27 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | 3.0 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$
IC342 XMM2 | 13.47 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ |  $<$2.8 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$
IC342 XMM3 | 5.47 | 2.75 | 2.75 | 5 | 10 | B0.5 | 2.0 | 7.0 | 33 | 5 | O9 | 7.1 | 40 | 5 | O9 | 7.0 | 33 | 5 | O9 | 7.1
N2403 XMM1 | 1.30 | 0$-$0.5 | 0.20 | 151 | $\cdots$ | $\cdots$ | $\cdots$ | 3.7 | 103 | 100 | B4.5 | 4.3 | 361 | 5 | O9 | 4.5 | 132 | 400 | A0 | 5.4
Ho II XMM1 | 0.87 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | 4.5 | 83 | $\cdots$ | $\cdots$ | 5.1 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | 83 | $\cdots$ | $\cdots$ | 5.1
Ho IX XMM1 | 1.18 | $\cdots$ | $\cdots$ | $\cdots$ |  $<$400 | A0 |  $<$1.0 | 3.2 | 25 | $\cdots$ | $\cdots$ | 3.8 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | 25 | $\cdots$ | $\cdots$ | 3.8
N4395 XMM1 | 1.13 | 0$-$0.5 | 0.30 | 19 | 5 | B4.5 | 0.5 | 2.8 | 28 | 50 | B3.5 | 4.1 | 443 | 5 | O9 | 4.2 | 240 | 10 | B0.5 | 4.3
N4490 XMM1 | 3.27 | $\cdots$ | $\cdots$ |  $<$ 180 | 30 | B2 | 1.0 | 5.0 | 171 |  $<$ 180 | B6.5 |  $<$ 5.5 | 418 |  $<$ 50 | B3.5 |  $<$ 5.3 | 364 |  $<$ 150 | B5.5 |  $<$ 5.5
N4490 XMM2 | 2.48 | $\cdots$ | $\cdots$ | $\cdots$ | 12 | B0.5 | 0.5 | 4.8 | 116 | $\cdots$ | $\cdots$ | 6.2 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | 116 | $\cdots$ | $\cdots$ | 6.2
N4490 XMM3 | 7.33 | $\cdots$ | $\cdots$ |  $<$ 75 | 3 | O6 | 2.5 | 4.9 | 146 |  $<$ 75 | B3.5 | 5.5 | 146 |  $<$ 75 | B.3.5 | 5.5 | 146 |  $<$ 75 | B.3.5 | 5.5
N4490 XMM4 | 5.75 | 0$-$1.5 | 0.88 | 5 | 10 | B0.5 | 0.5 | 4.4 | 133 | 15 | B1.5 | 4.9 | 192 | 1 | O5 | 4.6 | 507 | 4 | O8.5 | 5.6
N4490 XMM5 | 2.20 | 0.5$-$1.0 | 0.77 | 2 | 12 | B0.5 | 0.5 | 4.3 | 87 |  $<$ 50 | $\cdots$ |  $<$ 5.1 | 452 | 1 | O5 | 4.5 | 647 | 1 | O5 | 5.2
N4631 XMM1 | 1.69 | 0.50 | 0.50 | 7 | 10 | B0.5 | 0.5 | 4.7 | 69 | 5 | O9 | 5.3 | 235 | 5 | O9 | 5.0 | 671 | 5 | O9 | 5.3
N4736 XMM1 | 3.55 | 2.00 | 2.00 | 8 | 15 | B1.5 | 2.0 | 5.8 | 246 | 3 | O6 | 6.0 | 246 | 3 | O6 | 6.0 | 274 | 10 | B0.5 | 7.0
N5204 XMM1 | 0.50 | 0$-$0.5 | 0.15 | 20 | 15 | B1.5 | 0.0 | 3.9 | 52 | 30 | B2 | 4.7 | 309 | 3 | O6 | 3.9 | 321 | 18 | B1.5 | 4.8
M51 XMM2 | 0.73 | 0.5$-$1.0 | 0.75 | 12 | $\cdots$ | $\cdots$ | $\cdots$ | 3.9 | 94 | 20 | B1.5 | 4.7 | 453 | 3 | O6 | 3.7 | 354 | $\cdots$ | $\cdots$ | 4.7
M51 XMM6 | 4.62 | 0.5$-$1.25 | 0.75 | 48 | $\cdots$ | $\cdots$ | $\cdots$ | 4.6 | 157 | 5 | O9 | 4.7 | 157 | 5 | 87.1 | 4.7 | 229 | 150 | B5.5 | 5.8
M101 XMM1 | 0.12 | 0$-$0.75 | 0.42 | 11 | $\cdots$ | $\cdots$ | $\cdots$ | 3.5 | 74 | 15 | B1.5 | 4.0 | 484 | 5 | O9 | 4.6 | 708 | 10 | B0.5 | 5.7
M101 XMM2 | 0.90 | 0$-$1.0 | 0.38 | 56 | 3 | O6 | 1.0 | 2.8 | 84 | 5 | O9 | 4.0 | 812 | 3 | O6 | 3.3 | 178 | 300 | B8 | 4.6
M101 XMM3 | 1.12 | 0$-$0.5 | 0.35 | 206 | $\cdots$ | $\cdots$ | $\cdots$ | 3.5 | 678 | 3 | O6 | 3.8 | 678 | 3 | O6 | 3.8 | 930 | 1000 | F0 | 7.7
N5408 XMM1 | 0.51 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | 4.6 | 106 | $\cdots$ | $\cdots$ | 5.2 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ | 401 | $\cdots$ | $\cdots$ | 6.1
N6946 XMM1 | 1.87 | 1.00 | 1.00 | 12 | $\cdots$ | $\cdots$ | $\cdots$ | 3.4 | 30 | $\cdots$ | $\cdots$ | 3.9 | 225 | 12 | B0.5 | 4.6 | 225 | 12 | B0.5 | 4.6

Note. — (2) Extinction estimated using the column densities obtained by WMR
from X-ray spectral fits; (3) Range of extinction values for star-forming
regions around each ULX; (4) - (5) Average extinction and cluster age for the
star-forming regions around each ULX; (6) - (9) Results obtained for the 100
pc regions by comparing the OM colors with Starburst99 tracks: cluster age,
the earliest stars that survive in the cluster, given its age, and the mass of
the cluster; (10) - (13) Same results for the closest SFR for each ULX, plus
the offset distance to the center of the SFR. Some ULXs are actually located
inside these SFR, the offset distances for these are shown in bold. (14) -
(17) Same results for the youngest SFR around each ULX, plus the offset
distance to the center of the SFR. (18) - (21) Same results for the closest
SFR for each ULX, plus the offset distance to the center of the SFR. The final
row shows median values for each column.

Figure 1:  Color coded and adaptively smoothed OM UV images of NGC 2403 and
NGC 5204, showing the extraction regions of ULXs and nearby star forming
regions. The extraction region around the ULX position is 100pc for all ULXs,
but the extraction region around the surrounding star forming regions varies
in size based on the extent of the emitting region (see Section 2 for
details). The “X” marks the ULX position. Colors are for the three UV filters
on OM: UVW1 - red, UVM2 - green and UVW2 - blue.  Figure 2:  Magnitude offsets
caused by using the same sized aperture for each region across all OM filters.
The black lines are U$-$UVW1 color offsets versus aperture size, for point
sources with various count rates. The count rates are shown above each line.
The vertical line shows the radius of the aperture we used for many of our
sources as an example. For faint sources we used a maximum offset U$-$UVW1
$=-$0.15. For brighter sources we interpolated this value down to the minimum
U$-$UVW1 $=-$0.06. The B$-$UVW1 color offsets are shown (in red) and are very
similar to U$-$UVW1. Finally, for UVW1$-$UVW2 (in blue), the offsets are
smaller but still significant. In this case we used a maximum offset of
UVW1$-$UVW2 $=$ 0.06 for faint sources and zero offset for the brightest
sources.’  Figure 3:  The UVW1 magnitudes measured in 100 pc regions centered
on ULXs are plotted against the bolometric luminosities estimated from the
X-ray observations in WMR. Possible single star detections as defined in Table
The Local Environment of Ultraluminous X-ray Sources Viewed by XMM Newton’s
Optical Monitor are shown as circles while stellar associations are shown as
diamonds. Only four of these we estimate to have single stars as UV
counterparts (see details in Section 3.1), and these are labeled. We also show
with triangles the upper limits for ULXs with no detection (S/N $<$ 3). Most
UVW1 errors are smaller than the size of the symbols. The straight line shows
theoretical values for ULXs assuming the UVW1 fluxes are produced by a
standard accretion disk accreting at the Eddington limit. The dotted line
shows the theoretical irradiated disk emission. The marks on the top
horizontal axis show BH masses corresponding to the Eddington luminosities
below. Similarly, on the right vertical axis we mark star types corresponding
to the UVW1 magnitudes on the left axis.

Figure 4:  Color-color plots for the 100 pc regions in OM filters. The plotted
lines correspond to solar metallicity dwarfs and supergiants from the Kurucz
library. The dotted lines are reddened models (A${}_{V}=$ 1) for white dwarfs.
The short orange lines were calculated for ULXs with BH masses ranging from 5
to 10000 M⊙, assuming a standard accretion disk and self-irradiation (see text
for details). The arrow at the top left corner marks the maximum magnitude and
direction of the aperture correction.

Figure 5:  Color-color plots for the 100 pc regions, the symbols and tracks
are the same as in Figure 4. The arrow at the top left corner marks the
maximum magnitude and direction of the aperture correction

Figure 6:  Color-color plots for the 100 pc regions, the symbols and tracks
are the same as in Figure 4.

Figure 7:  Color-color plots for the 100 pc regions and nearby SFRs observed
with OM filters. The plotted tracks are Starburst99 instantaneous models with
solar metallicity. The data points for each track represent ages ranging from
1 to 900 Myr. We also plot reddened tracks for A${}_{V}=$ 1 (dotted) and
A${}_{V}=$ 2 (dashed), besides tracks with no extiction A${}_{V}=$ 0 (solid).
The arrow at the top left corner marks the maximum magnitude and direction of
the aperture correction. Finally we plot an average error bar over all data
points in the right bottom corner of each plot.

Figure 8:  Color-color plots, the Starburst99 tracks are the same as in Figure
7. We plot an average error bar over all data points in the right bottom
corner of each plot. The arrow at the top left corner marks the maximum
magnitude and direction of the aperture correction

Figure 9:  Color-color plots, the Starburst99 tracks are the same as in Figure
7. We plot an average error bar over all data points in the right bottom
corner of each plot.  Figure 10:  The UVW1 magnitude evolution of the
Starburst99 model with a cluster mass of 106 M⊙. This was used to estimate
masses for the clusters with reasonable age and extinction estimates, by
scaling from the Starburst99 model.

Figure 11:  Diagnostic diagrams for the runaway binary scenario. For the data
points we used the distances and star type estimates for the closest SFR
(Table 4, cols. 9 and 11). Left: the tracks correspond to different BH masses.
In some cases the ULXs appear to be located inside the SFRs (given the
resolution of the OM). These are plotted as circles. The solid lines
correspond to case B mass transfer, and the dashed lines for the case A. The
horizontal lines at the bottom of the plot show the kick velocities for each
BH mass. Right: Here the tracks correspond to different donor masses. On the
x-axis we plot both luminosities (bottom) and the corresponding BH masses
(top), assuming accretion at 10 times the Eddington limit. Here all the lines
are for case A mass transfer. The solid lines correspond to super-Eddington
accretion, while the dashed lines for accretion at the Eddington limit. We
plot with squares the ULXs that are associated with massive clusters ($>$105
M⊙). The black thick tracks mark the limiting kick velocity of 10 km s-1, for
the Eddington limit (dashed) and for super-Edington accretion (continuous).

Figure 12:  Diagnostic diagrams for the runaway IMBH scenario. These plots are
similar to Fig. 11, but now we assume that the ULXs contain IMBHs, and the
kicks are due to three-body interaction inside clusters. As in Fig. 11, the
ULXs that appear to be located inside SFRs are shown as circles. Left: as in
Fig. 11, the solid lines correspond to case B mass transfer, and the dashed
lines correspond to case A, but now we assume accretion at 0.1 the Eddington
limit. As in Fig. 11, we show ejection velocities on the horizontal lines at
the bottom. Right: the solid lines correspond to accretion at 0.1 the
Eddington limit, while the dashed lines for accretion at the Eddington limit.
The BH masses marked on the top x-axis correspond to the 0.1 Eddington
accretion case. As in Fig. 11, the black thick tracks mark the 10 km s-1 kick
velocity limit