# A Search for Gas-Rich Dwarf Galaxies in the Local Universe with ALFALFA
and the WIYN One Degree Imager

Katherine L. Rhode Department of Astronomy, Indiana University, 727 East Third
Street, Bloomington, IN 47405, USA<EMAIL_ADDRESS>Nicholas J. Smith
Department of Astronomy, Indiana University, 727 East Third Street,
Bloomington, IN 47405, USA  William F. Janesh Department of Astronomy, Case
Western Reserve University, 10900 Euclid Avenue, Cleveland, OH 44106, USA
John J. Salzer Department of Astronomy, Indiana University, 727 East Third
Street, Bloomington, IN 47405, USA  Elizabeth A. K. Adams ASTRON, Netherlands
Institute for Radio Astronomy, Oude Hoogeveensedijk 4, 7991 PD Dwingeloo, The
Netherlands  Kapteyn Astronomical Institute, University of Groningen,
Landleven 12, 9747 AD, Groningen, The Netherlands  Martha P. Haynes Cornell
Center for Astrophysics and Planetary Science, Space Sciences Building,
Cornell University, Ithaca, NY 14853, USA  Steven Janowiecki University of
Texas, Hobby-Eberly Telescope, McDonald Observatory, TX 79734, USA  John M.
Cannon Department of Physics & Astronomy, Macalester College, 1600 Grand
Avenue, Saint Paul, MN 55105

###### Abstract

We present results from an optical search for Local Group dwarf galaxy
candidates associated with the Ultra-Compact High Velocity Clouds (UCHVCs)
discovered by the ALFALFA neutral hydrogen survey. The ALFALFA UCHVCs are
isolated, compact HI clouds with projected sizes, velocities, and estimated HI
masses that suggest they may be nearby dwarf galaxies, but that have no clear
counterpart in existing optical survey data. We observed 26 UCHVCs with the
WIYN 3.5-m telescope and One Degree Imager (ODI) in two broadband filters and
searched the images for resolved stars with properties that match those of
stars in typical dwarf galaxies at distances $<$2.5 Mpc. We identify one
promising dwarf galaxy candidate at a distance of $\sim$570 kpc associated
with the UCHVC AGC 268071, and five other candidates that may deserve
additional follow-up. We carry out a detailed analysis of ODI imaging of a
UCHVC that is close in both projected distance and radial velocity to the
outer-halo Milky Way globular cluster Pal 3. We also use our improved
detection methods to reanalyze images of five UCHVCs that were found to have
possible optical counterparts during the first phase of the project, and
confirm the detection of a possible stellar counterpart to the UCHVC AGC
249525 at an estimated distance of $\sim$2 Mpc. We compare the optical and HI
properties of the dwarf galaxy candidates to the results from recent
theoretical simulations that model satellite galaxy populations in group
environments, as well as to the observed properties of galaxies in and around
the Local Group.

††facilities: WIYN (pODI and ODI), Arecibo (ALFA), WSRT††software: ODI-PPA
(Gopu et al., 2014), QuickReduce pipeline (Kotulla, 2014), IRAF (Tody, 1986,
1993), Python (https://www.python.org), Astropy (Astropy Collaboration et al.,
2013, 2018)

##  1 Introduction

The Arecibo Legacy Fast ALFA survey (ALFALFA; Giovanelli et al. 2005; Haynes
et al. 2018) was a survey for H i emission carried out from 2005$-$2012 with
the Arecibo L-band Feed Array (ALFA) instrument on the 305-m Arecibo radio
telescope. The survey covered $\sim$7000 sq. degrees of sky in drift-scan
mode, detecting sources of 21-cm line emission without regard to the locations
of optical sources. The final catalog produced by the ALFALFA survey includes
$\sim$31,500 H i sources out to a redshift of $z\sim 0.06$, a large fraction
of which had never before been observed via their 21-cm emission (Haynes et
al., 2018).

Because of its excellent sensitivity, ALFALFA was able to detect objects with
only modest amounts of H i ($\sim$$10^{4}-10^{5}$ $M_{\odot}$) out to
distances beyond the edge of the Local Group. A new class of potentially
nearby, low-mass H i sources was identified in the ALFALFA survey data and
dubbed Ultra-Compact High Velocity Clouds (UCHVCs; Giovanelli et al. (2010)).
Roughly 100 of these objects were found in ALFALFA; Adams et al. (2013) laid
out a set of criteria for sources that qualify as UCHVCs and presented an
initial catalog. These objects are compact (angular diameter $<$30′), low-mass
(H i masses of $10^{5}-10^{6}$ $M_{\odot}$ at 1 Mpc), isolated from other
sources in terms of both sky position and velocity, and have velocities that
help distinguish them from Galactic High Velocity Clouds while making them
likely to be located within the Local Volume. Perhaps most importantly,
ALFALFA sources are classified as UCHVCs only if they lack an obvious optical
counterpart when matched with existing optical survey data from the Sloan
Digital Sky Survey (SDSS; York et al. 2000; Eisenstein et al. 2011) or the
Digitized Sky Survey (DSS).

The properties of the UCHVCs imply that these objects could represent a
population of low-mass, gas-bearing dark matter halos located in and around
the Local Group; this is the “minihalo hypothesis” described in Giovanelli et
al. (2010) and Adams et al. (2013). Accordingly, our group has been observing
a selection of the UCHVCs at optical wavelengths in order to investigate their
nature and the possibility that the UCHVCs are nearby gas-rich dwarf galaxies
that have previously gone undetected. We image the UCHVC locations at optical
wavelengths with the WIYN 3.5-m Observatory111The WIYN Observatory is a joint
facility of the NSF’s National Optical-Infrared Astronomy Research Laboratory,
Indiana University, the University of Wisconsin-Madison, Pennsylvania State
University, and Purdue University. and systematically search the area around
the location of the H i detection to identify any stellar populations that may
be associated with the neutral gas. The first UCHVC we observed with WIYN led
to the discovery of the nearby gas-rich dwarf galaxy Leo P (Giovanelli et al.,
2013; Rhode et al., 2013) which is located just outside the Local Group at
$\sim$1.6 Mpc, is extremely metal-deficient, and is the lowest-mass galaxy
known that is actively forming stars (Skillman et al., 2013; McQuinn et al.,
2015).

After the discovery of Leo P, we initiated a systematic optical observing
campaign that targeted UCHVCs from the catalog presented in Adams et al.
(2013) as well as additional sources drawn from the low-velocity ALFALFA
grids. We used wide-field imaging in two broadband filters, resolved stellar
photometry, and a color magnitude diagram (CMD) filtering technique to search
for the presence of possible stellar populations associated with the HI clouds
over a distance range of 250 kpc to 2.5 Mpc. The first part of the campaign
was carried out with the WIYN One Degree Imager with a partially-filled focal
plane – an instrument referred to as pODI – and focused mainly on objects with
higher H i column density values. In addition, we used the Very Large Array
(VLA) and the Westerbork Synthesis Radio Telescope (WSRT) to acquire and
analyze H i synthesis data for several of the targets in the pODI sample, in
order to obtain more detailed information about the objects’ H i distributions
and kinematics (e.g., Adams et al., 2015, 2016; Paine et al., 2020; Bralts-
Kelly et al., 2020). Optical imaging studies were also carried out by other
groups to search for counterparts to the UCHVCs and to other potentially
nearby gas-rich sources found in surveys like GALFA-HI (Bellazzini et al.,
2015a, b; Sand et al., 2015); the results from these studies will be discussed
further in later sections of this paper.

A total of 23 UCHVCs that were observed with WIYN pODI were analyzed in the
first phase of the project. In Janesh et al. (2015), we described the WIYN
survey and methods and presented some initial results for AGC 198606, a UCHVC
that is close in both projected spatial location and velocity to the Local
Group dwarf galaxy Leo T (Irwin et al., 2007; Ryan-Weber et al., 2008).
Results for another UCHVC, AGC 249525, were initially presented in Janesh et
al. (2017) and the complete set of results for the first 23 UCHVCs we analyzed
was published in Janesh et al. (2019) (hereafter J19). J19 identified a total
of five UCHVCs that had possible faint stellar counterparts in the pODI
imaging data. These five counterparts had estimated distances ranging from
$\sim$350 kpc to 1.6 Mpc and estimated total optical magnitudes between
$M_{V}$ $=$ $-$1.4 and $-$7.1 mag.

The second phase of our UCHVC follow-up imaging campaign was carried out with
the upgraded One Degree Imager (ODI) camera on WIYN, which was commissioned in
2015 and which provides a field-of-view (FOV) more than three times larger
than that of pODI. In a series of observing runs between 2015 and 2021, we
acquired deep WIYN ODI images of more than 30 additional UCHVC targets in two
broadband filters (SDSS $g$ and $i$). During this phase of the project, we
made improvements to our detection and analysis methods and then applied the
updated methods to the ODI sample of objects. We also used the improved
methods to reanalyze the imaging data of the five UCHVCs in the pODI sample
for which J19 identified potential stellar counterparts.

In this paper, we present the results from the second and final phase of our
campaign to obtain follow-up optical imaging of ALFALFA UCHVCs. In Section 2,
we describe the selection criteria and properties of the UCHVCs we observed
and analyzed. Section 3 summarizes the procedures we used to acquire imaging
observations and produce source catalogs and calibrated broadband optical
photometry for each UCHVC field. In Section 4, we describe the steps we
carried out to search for stellar counterparts to the HI sources and to assess
the significance and validity of the dwarf galaxy candidates that we have
identified. We also discuss the improvements that were made to our analysis
methods in the second phase of the project. Section 5 presents the full
results of this analysis. In the last section of the paper, we examine our
results in the context of recent relevant theoretical work and other
observational studies aimed at detecting and studying gas-rich dwarf galaxies
in and around the Local Group.

##  2  The Sample of UCHVCs Analyzed for This Work

The UCHVCs we observed were drawn from the ALFALFA survey data (Haynes et al.,
2018) using criteria laid out in detail in Adams et al. (2013). The primary
criteria that were used to select sources are as follows:

  * •

The source must have an H i major axis, $a$, less than 30$\arcmin$ in order to
have a size consistent with expectations for the baryonic component of low-
mass dark matter halos. This corresponds to a diameter of $\sim$2 kpc at a
distance of 250 kpc, $\sim$9 kpc at 1 Mpc, and $\sim$20 kpc at 2.5 Mpc.

  * •

The source must have an ALFALFA signal-to-noise ratio greater than 8, as
defined in Equation 2 in the ALFALFA 40% catalog (Haynes et al., 2011).

  * •

The source must have no more than three neighboring HI structures within 3
degrees on the sky or 15 km s-1 in velocity space in order to ensure that it
is sufficiently isolated. It must also have at least a 15-degree separation
from any previously known high-velocity cloud (HVC) complexes.

  * •

The source must have a heliocentric velocity between $-$500 and 1000 km s-1,
so that it is likely to reside in the Local Volume.

  * •

The source must have an absolute velocity with respect to the Local Standard
of Rest (LSR) greater than 120 km s-1, to avoid contamination with Galactic
HVCs. This criterion was in some cases relaxed when all other criteria were
met, since some nearby dwarf galaxies have velocities lower than this value
(e.g., Leo T, with $v_{\rm LSR}$ $\sim$60 km s-1; Irwin et al., 2007; Ryan-
Weber et al., 2008).

  * •

The source must have no clear optical counterparts in either SDSS or DSS.

Approximately 100 ALFALFA sources were identified as UCHVCs, and we selected a
sample of roughly 60 sources from that original list that were best suited for
optical follow-up (e.g., strong HI detections, relatively high column
densities, and/or potentially interesting for other reasons, such as sky
position or velocity); see J19 for more details about how the optical follow-
up sample was selected. As mentioned in the Introduction, 23 of these objects
were analyzed in the first phase of the project and presented in J19. We
observed 35 more objects after the pODI instrument was replaced with ODI. The
substantially larger FOV of ODI (see Section 3) necessitated modifications to
our processing and detection steps, so it made sense to split the project into
two phases in this way.

The final sample of objects being presented here includes 25 objects that were
observed with ODI and for which we were able to obtain a complete set of
observations, i.e., high-quality, deep imaging in both filters ($g$ and $i$).
In addition, we include one more UCHVC (AGC 227977) that already had existing
pODI imaging data, but that we had planned to re-observe with ODI under better
sky conditions. We were not able to acquire a complete set of high-quality ODI
images for this object, but in the end we deemed the existing pODI images to
be of sufficient quality that we could include the object in the current
sample. Nine more UCHVCs were observed, but the imaging data could not be
fully analyzed for a range of reasons (e.g., observed under poor sky
conditions, good-quality images obtained in only one filter); the objects are
mentioned here for completeness.

The ALFALFA-derived properties of the 26 UCHVCs we analyzed for the second
phase of the project are presented in Table 1. The table includes the source
designation, sky position, H i line flux ($S21$), recessional velocity ($cz$),
H i line FWHM ($W_{50}$), mean angular diameter of the H i detection ($a$),
mean H i column density (log ${\rm N_{HI}}$), H i mass at an assumed distance
of 1 Mpc, the dates of the observations, and the detection status from our
analysis.

Information about the nine other UCHVCs that could not be analyzed as part of
this phase of the project are presented in Table 2. The columns are the same
as those shown in Table 1 except that no detection status is given since a
full analysis could not be done. We include information about the observations
we were able to acquire when applicable. The objects that fall into this
category typically had poor-quality data and/or imaging in only one filter.

Table 1 shows that the objects we analyzed possess a fairly broad range of H i
properties. The objects in our sample have H i flux values that range from
0.49 to 7.95 Jy km s-1, with most sources between 0.6 to 2.0 Jy km s-1. In
general, we prioritized H i sources with higher column densities during the
data acquisition phase of the project, under the assumption that these targets
would be more likely to host a dwarf galaxy; one (albeit indirect) consequence
of this is that the sources we were unable to analyze because of insufficient
data (Table 2) tend to have smaller total HI fluxes than the objects in the
sample shown in Table 1. The H i major axes of our objects range from
5.3$\arcmin$ to 17.2$\arcmin$ (approximately 1.5 to 5.0 kpc at a distance of 1
Mpc) and they span a recessional velocity range from $-$452 to $+$320 km s-1.

##  3 Observations & Initial Data Reduction

We obtained optical imaging of the UCHVC sources with the WIYN 3.5-m telescope
at Kitt Peak National Observatory over the course of several observing
seasons; the year and semester of the observations of each source are given in
Table 1. Observations carried out prior to 2015 were taken with pODI, which
had a central imaging area that consisted of a 3 x 3 array of orthogonal
transfer arrays (OTAs). Each individual OTA is comprised of an 8 x 8 grid of
CCD detectors. The pODI configuration provided a 24′ by 24′ FOV and a pixel
scale of 0.11″ pixel-1. The upgraded ODI camera includes a 5 x 6 arrangement
of OTAs, yielding a 40′ x 48′ FOV and the same pixel scale as pODI.

Each of the UCHVCs was imaged with a series of nine 300-second exposures in
both the SDSS g and i filters, and the telescope was dithered slightly between
exposures in order to fill in the gaps between the CCD detectors and the OTAs;
this yielded a total integration time of 45 minutes in each filter. The images
were transferred to the One Degree Imager Pipeline, Portal, and Archive (ODI-
PPA) system222The ODI Portal, Pipeline, and Archive (ODI-PPA) system is a
joint development project of the WIYN Consortium, Inc., in partnership with
Indiana University’s Pervasive Technology Institute (PTI) and NSF’s NOIRLab.
(Gopu et al., 2014; Young et al., 2014) for storage and processing. The
QuickReduce data reduction pipeline (Kotulla, 2014) was used to carry out
pixel masking, crosstalk and persistence correction, overscan subtraction,
bias and dark subtraction, flat-field correction, pupil ghost correction, and
cosmic ray removal on the images.

After the QuickReduce processing, the images were further reduced and stacked
using the odi-tools software routines, a suite of Python routines developed by
William Janesh and Owen Boberg to help facilitate the analysis of the J19
UCHVCs sample and other pODI and ODI data (Janesh, 2018). The images were
illumination-corrected and reprojected to the same pixel scale and coordinate
reference frame. Sets of images taken in the same filter were then scaled to a
common flux level and combined into a single science-ready image. The combined
images in the two filters were aligned and then trimmed to 20000 pixels by
22000 pixels ($\sim$37′ x 40′) to cut out the regions without full coverage in
the dither pattern. The mean full-width at half-maximum of the point-spread
function (FWHM PSF) in the final combined images was typically
${}_{<}\atop{{}^{\sim}}$1.0″ in both filters. In cases where we observed an
object on multiple nights with varying sky conditions, we used the images with
the best seeing values to construct the final combined image in a given
filter.

All of the UCHVC targets fall within the footprint of SDSS, so we measured the
magnitudes and colors of several hundred SDSS stars that appeared within each
of the ODI frames in order to derive photometric calibration coefficients
(zero points and color terms) that we could then use to calculate calibrated
magnitudes and colors for the other sources in the final combined images.
Typical uncertainties on the photometric zero points in the calibrated $g$ and
$i$ magnitude equations are $\sim$0.02 mag.

##  4 Searching for Stellar Populations Associated with the HI Clouds

###  4.1 Applying a Color-Magnitude Diagram Filter to the Stellar Catalogs

When we analyze our optical images of the UCHVCs, the goal is to search for
collections of stars – which we refer to as “stellar overdensities” – with
properties like those of the stellar populations one would expect to be
associated with a Local Group dwarf galaxy. We have developed a largely
automated procedure to carry out this search as well as to characterize the
statistical significance of any stellar overdensities we find. The overall
method we use was laid out in Janesh et al. (2015); as mentioned earlier, we
have since made some improvements to the method, which are described in this
section. Additional details can be found in Smith (2022).

We use a processing pipeline that calls Python routines and IRAF tasks and
includes both automated steps and steps that require input from the user. The
main steps in the process are as follows. First, we carry out detection and
aperture photometry of the sources in the WIYN images that are above a certain
count threshold – typically, 3.5 to 4 times the standard deviation of the sky
background level in the images. The source lists generated from the two images
are then matched, and extended objects are eliminated based on the difference
in magnitudes measured in two different apertures. We measure the instrumental
magnitude of each source using a small aperture with a radius equal to the
average FWHM of the image (magnitude $m1$) and a larger aperture with a radius
equal to twice that value (magnitude $m2$). Bright point sources fall along a
clearly-defined sequence when plotted in the $m2$ vs. $m1-m2$ plane; we use
that sequence to define the center of our selection region and then select
sources within 2.5$\sigma$ of that center. The selection region is thus wider
at fainter magnitudes where the photometric uncertainties are larger; see
Janesh (2018) and Smith (2022) for more details. Aperture photometry of the
point sources that survive the extended source cut is then carried out in both
images, using an aperture radius equal to the average FWHM PSF of the image.
An aperture correction is computed from photometric measurements of selected
bright stars across the image and then applied to the instrumental magnitudes
of all of the point sources. Final calibrated magnitudes and colors are
calculated for each point source in the field by applying the appropriate
photometric calibration coefficients along with Galactic extinction
corrections calculated from the Schlegel et al. (1998) dust maps and the
relations in Schlafly & Finkbeiner (2011). Typical 5-$\sigma$ detection limits
in the images are $g$$\sim$25.4 mag and $i$$\sim$24.4 mag.

Next we construct a color-magnitude diagram (CMD) filter that can be applied
to the final photometry catalog for a given field. The construction of the
filter was described in detail in Janesh et al. (2015) and followed the basic
method outlined in Walsh et al. (2009) using stellar isochrones from Girardi
et al. (2004). Because UCHVCs are expected to host old, metal-poor stars, the
CMD filter is designed to select stars with ages between 8 and 14 Gyr and
metallicities between Z = 0.0001 and 0.0004. In Janesh et al. (2015), we used
our WIYN observations of Leo P to test the CMD filtering method and
demonstrated that the filter and our detection methods are effective at
finding a stellar population that consists of only a modest population of red
giant branch (RGB) stars. The position and extent of the filter on the CMD is
defined by the locations of the isochrones for the chosen age and metallicity
ranges, as well as by the distance to the putative stellar population. We
shift the filter in the vertical (magnitude) direction on the CMD over a range
of magnitudes in increments of 0.01 mag, to cover a corresponding distance
range of 250 kpc to 2.5 Mpc. Applying the CMD filter at each distance step to
our photometric catalog yields a subset of stars in a given UCHVC field that,
given their $g$ and $i$ photometric measurements and associated uncertainties,
lie in a region of the CMD that is consistent with the location and boundaries
of the filter. Figure 1 shows an example of a CMD filter being applied to our
data; the black dots show the CMD positions of the point sources that appear
in one of the UCHVC fields, the blue solid line marks the location of the CMD
filter corresponding to a distance of 400 kpc and the age and metallicity
ranges mentioned above, and the red dots are the sources that are selected by
the CMD filter in this case.

Figure 1: An example of the CMD filter that is applied to the photometric
catalog derived from the WIYN ODI imaging of each UCHVC field in order to
select stars with properties expected for a dwarf galaxy at a given distance
(see Sec. 4.1). Black points are point sources located within the images and
red points are those that coincide with the location of the filter given their
$g$ and $i$ magnitudes and associated 1-$\sigma$ uncertainties. Error bars
showing the typical photometric uncertainties at various magnitudes are
plotted in blue on the right-hand side. The filter shown in this example
corresponds to a distance of 400 kpc. Figure 2: An example of the effect of
the cross-matching step that is carried out on source catalogs of fields in
which potential stellar overdensities have been identified. Sources in the
field are cross-matched with Gaia and SDSS catalogs in order to identify and
remove contaminating foreground objects (Galactic stars) and background
objects (galaxies and AGN) from our star lists. Sources that are retained are
shown with black points and sources that are removed are shown with red
points. In this example field, approximately one-third of the objects in our
original source list were removed as likely Galactic stars and $\sim$10% were
removed as likely background galaxies or AGN; these percentages are fairly
typical, although the exact proportions of foreground and background objects
and remaining point sources depend on quantities like the Galactic latitude
and exact detection limits of each set of images.

###  4.2 Searching for Overdensities in the Spatial Distribution of CMD-
Filtered Stars

Once we have identified the set of stars selected by a given CMD filter, we
search for clustering among the stars to identify any stellar populations that
might be associated with the H i source. We begin by smoothing the spatial
distribution of the stars with a Gaussian kernel. We use smoothing kernels
with a diameter of both 2$\arcmin$ and 3$\arcmin$ in order to mimic the range
of typical angular sizes that a dwarf galaxy might have if it were located
within the distance range being probed. We then estimate the density of stars
at each location in the smoothed image as a function of the mean density of
the CMD-filtered stars across the image. Any region of the image where the
stellar density exceeds the mean value represents a potential stellar
overdensity associated with the H i source.

In order to determine the significance of a given overdensity (i.e., how
likely it is to be a genuine overdensity rather than a result of random
fluctuations in the point source spatial distribution), we carry out a series
of 10,000 Monte Carlo experiments corresponding to each CMD filter position.
We distribute the same number of CMD-filtered stars at random locations across
the field and repeat the smoothing steps. Any overdensity identified in the
real data can then be compared to those found in the randomly-generated data.
The significance of the overdensity in the real data is quantified by
calculating the number of overdensities in the random realizations that are
less dense than the real overdensity. After carrying out these steps for every
distance and for both smoothing radius sizes, we can compare the overall
results across the entire parameter space that has been searched and look for
the most significant detections that appear within a given UCHVC field.

###  4.3 Assessing the Detected Stellar Overdensities and Cross-Matching with
the Gaia and SDSS Catalogs

The next stage of the analysis includes a variety of steps that are applied as
appropriate, depending on the results of the CMD filtering and smoothing
process. The first step in all cases is to examine the location and
significance of the overdensities revealed by the CMD filtering process at
every step of the distance range being sampled (250 kpc to 2.5 Mpc). We focus
on detections that have high significance (approximately 90% or higher) and
that are close in projection to the H i centroid position – usually within
8$\arcmin$, which translates to a separation of 2.3 kpc at a distance of 1
Mpc. We sometimes consider overdensities that have significance in the
80$-$90% range, or are up to 10$\arcmin$ away from the H i centroid, if their
other characteristics suggest they could be a genuine dwarf galaxy candidate.
If no overdensities are found that meet these criteria, we classify the UCHVC
target as a non-detection and do not carry out any further analysis on the
imaging data for that UCHVC. If multiple unique detections (i.e., at different
sky positions and/or distances) are found, each is assessed with the next
steps. We note also that the steps described here are carried out on the
results produced using both of the smoothing kernels (i.e., smoothing kernels
with diameters of 2$\arcmin$ and 3$\arcmin$), and we take the results derived
from both kernel sizes into account when selecting and classifying the various
overdensities that are identified.

Once we have identified possible stellar overdensities associated with the H i
sources, we examine the properties of the individual stars that are within
3$\arcmin$ of the center of the detected overdensity. We look at the CMD
locations of the stars that compose the detection to judge whether the stars
match the expected morphology of the RGB of a genuine dwarf galaxy at the
specific distance to which the filter location corresponds. We generally look
for a detected overdensity that has stars that populate a fair fraction of the
RGB region of the CMD filter, as well as stars that lie in the Horizontal
Branch (HB) region when applicable. We also compare the numbers and positions
of the stars in the detected overdensity to those of the stars in a reference
circle of similar size that is placed at a random location near the edges of
the image, far from the H i centroid position (to ensure that it does not
overlap with any detection circles). If the CMD of the stars in the detection
circle is similar in appearance and number of stars to the CMD of the
reference circle, then the detection is deemed less convincing.

We also inspect how closely the stars within the detection CMD are clustered
on the image. In some cases, the stars that comprise a given detected
overdensity may not actually be very close to each other spatially (e.g., the
detection might consist of several stars that are scattered around the edges
of a particular smoothed region that makes up the detected overdensity), which
makes the detection less convincing as a genuine dwarf galaxy candidate. In
addition, we examine the radial profiles of the individual stars within the
detection, and the objects in the region around the detection area, because in
some cases (for example) clusters of faint, distant unresolved background
galaxies can masquerade as collections of nearby stars with colors like RGB
stars. To help with this step, we also perform a quick cross-reference of the
stars within the over-density with source catalogs from the SDSS survey, to
see whether the sources in and around the detection are actually galaxies.
Finding a detected overdensity that is close (in projection) to an obvious
background galaxy group or cluster suggests that the detection may actually be
made up of unresolved background galaxies rather than being a genuine stellar
association in or near the Local Group.

After these steps are completed, the UCHVC targets with detected overdensities
are assigned into separate categories based on our initial assessment of the
likelihood that we may have detected the presence of a genuine dwarf galaxy.
UCHVCs with overdensities that no longer seem convincing after these checks
are designated as ”non-detections” (see Table 1) and are not analyzed further.

For the UCHVC targets that remain, we cross-match the point source catalogs
with the Gaia333This work has made use of data from the European Space Agency
(ESA) mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia
Data Processing and Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. EDR3 catalog (Gaia
Collaboration et al., 2016, 2021) and remove foreground stars (by rejecting
any source with a proper motion greater than three times the uncertainty on
that quantity) as well as known AGN (by removing sources that match those in
the list of AGN used to calibrate the Gaia EDR3 celestial reference frame;
Gaia Collaboration et al., 2021). We also make use of the star/galaxy
classification in the SDSS catalog to remove any sources that are likely to be
galaxies. Before implementing the latter step, we examined how reliable the
SDSS star/galaxy classification was by comparing it to the classification used
by the Cosmic Assembly Near-infrared Deep Extralactic Legacy Survey444This
work is based on observations taken by the CANDELS Multi-Cycle Treasury
Program with the NASA/ESA HST, which is operated by the Association of
Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.
(CANDELS; Grogin et al., 2011; Koekemoer et al., 2011; Stefanon et al., 2017).
CANDELS was carried out with the Hubble Space Telescope (HST) and therefore
genuine point sources can be distinguished from background galaxies more
effectively in CANDELS data than in data from a ground-based survey. We found
that for sources in common between SDSS and CANDELS, the CANDELS star/galaxy
classification matched the SDSS classification $>$94% of the time. Since mis-
identification occurred for only a small fraction of sources, and the number
of objects in our source lists that are identified as galaxies in SDSS is also
relatively small, we decided that applying the SDSS classification to our
source lists was warranted because it would help us weed out spurious dwarf
galaxy candidate detections. Figure 2 shows that the majority of sources that
are eliminated when we cross-match our point source catalogs with Gaia EDR3
and SDSS are bright (with $i$ $<$ about 20 mag) and are excluded because they
are Galactic foreground stars.

After removing the contaminants found in the catalog cross-matching, we repeat
the CMD filtering and smoothing steps, again assess each over-density that is
detected in a given UCHVC field, and update the initial classification that
was assigned. The final results of the search process are sorted into three
categories: ”non-detections” (UCHVCs that have no convincing evidence for an
optical counterpart); ”possible detections” (UCHVCs that may have an optical
counterpart that warrants further follow-up), and ”best detections” (UCHVCs
that have a convincing or likely optical counterpart).

For objects that fall within the latter two categories, we estimate the H i
and optical properties for the UCHVC and its counterpart at the relevant
distance. We take the distance of the object to be the CMD filter distance
that corresponds to the most significant detection that occurs for a given
overdensity. We also tabulate the range of distances over which a given
detection remains above some threshold for significance – 90% for
overdensities that have a peak significance $\geq$90%, and 80% for
overdensities with a peak significance between 80$-$90% – and use this to help
define an uncertainty on the estimated distance (see Sections 5 and 6).

###  4.4  Quantifying the Detection Completeness

As part of the analysis, we carry out artificial star tests to quantify the
photometric depth and detection limits of each set of UCHVC images. Before
running the full set of tests, we confirmed that completeness testing
performed on a 4000x4000-pixel subsection of the images yields the same
results as tests performed on the full-size ODI images. Since source crowding
is not an issue in our UCHVC target fields (they are empty of strongly
clustered sources, with the exception of a field that includes the globular
cluster Pal 3; see Sec. 5.1.2 for more discussion), we found that running
artificial star tests on portions of the images is a valid approach. The rest
of the steps were carried out on these smaller-sized images.

We characterized the point spread function (PSF) within each of the images by
fitting the light distributions of several dozen bright but unsaturated stars
within the frames. We then generated artificial stars and injected them into
the images, 400 stars at a time in steps of 0.2 mag, until we had spanned the
range of stellar magnitudes present in the real data. We followed our
detection procedures to recover and measure the magnitudes of the artificial
stars and used that information to create a completeness curve (which
quantifies the proportion of stars recovered as a function of magnitude) for
that image and filter. We created 50% completeness curves for the CMD of a
given UCHVC field by combining the completeness information in each of the
filters. Specifically, we constructed a grid of $i$ and $g-i$ values and then
computed the completeness at a given position in the grid by multiplying the
$i$ magnitude completeness by the appropriate completeness in the $g$ filter
based on the color at that grid position; for example, for the point at $i$
$=$ 25 mag and $g-i$ $=$ 1, we multiplied the detection completeness in $i$ at
25 mag by the completeness in $g$ at 26 mag. We carried out this calculation
over the entire range of source magnitudes and colors that appear in our CMDs
(i.e., $i$ $\geq$ 26 mag, $g-i$ from $-$1.5 to $+$4). We then interpolated
between the grid values to define the 50% completeness curve for a given UCHVC
field. These 50% completeness curves are included in the results presented in
Section 5.

###  4.5 Changes to the Original Detection Methods

Before analyzing the WIYN ODI data set, we examined each of the detection
steps carried out on the J19 sample to explore whether modifying the steps
might increase our chances of detecting dwarf galaxy candidates associated
with the UCHVCs. We made two substantial modifications to our process.

The first was a change in the CMD filtering step. As described in Sec. 4.1,
the boundaries of the CMD filter are defined by the set of isochrones that
characterizes the expected stellar population of a UCHVC. When we apply the
filter, we select sources that have $i$ magnitudes and $g-i$ colors that would
fall within the filter given the relevant photometric uncertainties. When we
examined this step in the detection process, we discovered a numerical error
in the original code that was applied to the J19 WIYN pODI data which made the
CMD filter less restrictive than intended. The effect of the error was more
pronounced at faint magnitudes, where the photometric errors are large, but
the impact was mitigated because the code also included a photometric error
cut that removed sources with $i$ magnitude errors $\geq$0.2 mag or $g-i$
color errors $\geq$0.28 mag. Nevertheless, the end result of this coding error
was that the pODI data set yielded a relatively large number of detections
that were later rejected as false when they were examined individually.

After correcting this error and then testing various options for how to apply
the CMD filter to the ODI imaging data, we decided to implement a relatively
restrictive source selection, allowing only those objects that are consistent
with the boundaries of the CMD filter given the 1-$\sigma$ photometric errors
of the object. We also implemented a change so that the code uses the
uncertainties in the $g$ magnitude and $i$ magnitude (instead of the
uncertainties in the $i$ magnitude and $g-i$ color, as was the case in the
original code) to determine whether a source should be accepted. More
specifically, we imposed the condition that sources were selected only if the
$g$ magnitude and $i$ magnitude errors were both $\leq$0.2 mag. This prevented
sources with, for example, small $i$ uncertainties but larger $g$
uncertainties from passing the filtering step. The final result is a more
restrictive CMD filter than the filter originally applied to the J19 pODI data
set; the revised filtering process only selects stars with reliable photometry
in both filters that match closely with the metallicity and age range of the
expected stellar population, which reduces the number of spurious detections
appearing in each of the UCHVC fields.

The second major change made to the detection process was to add the cross-
matching between our sources and objects in the Gaia and SDSS catalogs that
was described in Section 4.3. This added step serves to much more effectively
remove both foreground and background contaminants from each of the UCHVC
fields (see Fig. 2), which was especially helpful for processing the ODI data
set, given the large field of view and the large numbers of sources detected
in each field.

We used the improved series of steps to process all of the objects in the ODI
sample listed in Table 1. We also decided to reanalyze the WIYN pODI imaging
data for the five UCHVCs in the J19 sample that yielded possible dwarf galaxy
candidates. Results from the current sample of UCHVCs, as well as the re-
analysis of the five pODI objects from J19, are presented in the next section.

##  5 Results

###  5.1 Results for the Sample of UCHVCs Analyzed for This Work

We carried out the above-described processing and analysis steps for the 26
objects in the sample listed in Table 1. Twenty of the UCHVCs in this sample
had no stellar overdensity in the images that satisfied our criteria for a
dwarf galaxy candidate, and are thus classified as ”Non-Detections” and
labeled ”ND” in the table. The first 19 non-detections are discussed in Sec.
5.1.1, and one object in this category warrants a separate discussion in Sec.
5.1.2. Five objects classified as ”Possible Detections” are described in Sec.
5.1.3, and the one object classified as a ”Best Detection” is discussed in
Sec. 5.1.4.

####  5.1.1 Objects Classified as ”Non-Detections”

For most of the objects in this category, the classification was
straightforward because no stellar overdensities were found that had
significance in the $\sim$80$-$90% range or higher and were near the H i
centroid. In a few other cases, significant overdensities were identified
initially, but then not found to be believable upon close inspection of the
individual sources located within the overdensity and in the surrounding
region. One object (AGC 335755) made it through the first round of screening
to the final stages of analysis – i.e., the cross-matching with the Gaia and
SDSS catalogs – but was no longer deemed convincing after the contaminating
foreground and background objects were removed from the source list. In any
case, these 19 UCHVCs show no evidence that they host a dwarf galaxy candidate
within the distance range that we tested (250 kpc to 2.5 Mpc), given the
detection limits of our images. We note that other searches for optical
counterparts to compact H i clouds have detected dwarf galaxies at larger
distances ($\sim$3$-$20 Mpc; Bellazzini et al., 2015a, b; Sand et al., 2015;
Tollerud et al., 2015), so there is still the possibility that these UCHVCs
are associated with a stellar counterpart that is farther away than 2.5 Mpc
and therefore beyond the scope of our survey. On the other hand, our images of
the UCHVCs in this category contain no faint optical counterparts (either
resolved-star counterparts or unresolved faint optical emission) like the ones
detected in those studies.

####  5.1.2 An H i Cloud Near in a Globular Cluster?

One more UCHVC in the ”Non Detection” category requires a more detailed
discussion. AGC 501816 was included in the ODI sample because it lacked an
obvious optical counterpart and its other properties satisfy the criteria for
UCHVCs, except that its $|v_{LSR}|$ value, 99 km s-1, is below the nominal 120
km s-1 minimum requirement. In addition to being of interest as a possible
compact H i cloud, AGC 501816 is notable because of its proximity to the Milky
Way globular cluster (GC) Pal 3. The centroid of the H i source lies 7.3′ away
in projection from Pal 3 and the GC appears within our ODI images. This
projected angular separation translates to only 195 pc at the distance of Pal
3 (91.9 kpc; Hilker, 2006). Moreover, the GC and the UCHVC have similar
heliocentric velocities: Pal 3 has a $v_{\rm helio}$ of 94.0$\pm$0.8 km s-1,
only 13 km s-1 below that of AGC 501816 (Table 1).

Figure 3: Top: Diagram showing the positions of sources across the field-of-
view of our ODI images of AGC 501816, which includes the Galactic GC Pal 3.
Gray points mark the locations of objects in the final stellar catalog (after
cross-matching with the Gaia EDR3 and SDSS catalogs and removing
contaminants). The ALFALFA H i source is shown with a solid green ellipse. Pal
3 is the cluster of point sources $\sim$7 arcminutes away from the H i
centroid. The half-light radius and tidal radius of the GC from Baumgardt &
Hilker (2018) are marked, along with the 3 $r_{h}$ circle used to remove Pal 3
stars before we searched for stellar overdensities associated with the H i.
Two detected overdensities are marked with cyan and purple circles. The
stellar overdensities and the H i source are either inside or overlapping with
the tidal radius of Pal 3. Bottom: The CMDs for the two stellar overdensities
detected in the AGC 501816 field after all stars within 3 $r_{h}$ of Pal 3’s
center were removed. Gray points are all sources inside a 3′ radius around the
peak of the overdensity and red points are sources within that radius that
fall within the CMD filter at the given distance. For reference, we have over-
plotted an isochrone from Girardi et al. (2004) (solid blue line) with the age
and metallicity of Pal 3 (10 Gyr, [Fe/H] $=$ $-$1.7 or Z = 0.0004; Hilker
(2006)), shifted to match the distance at which each overdensity is detected
(80 kpc or 90 kpc).  Figure 4:  The UCHVC AGC 501816 is located on the sky
roughly 7$\arcmin$ from the outer Galactic halo globular cluster Pal 3. We ran
the detection pipeline on the images of the AGC 501816 field with Pal 3
present and with Pal 3 removed; this set of plots shows some of the results
from the searches carried out with the GC present in the images. The upper
left plot shows the locations of the stars detected in the field (gray
points), the stars selected by the CMD filter (red points), and the H i
ellipse for the UCHVC from ALFALFA (green ellipse). The region of the detected
overdensity is marked with a magenta circle of radius 3′. A yellow circle of
the same size is placed at a random location in the outskirts of the field and
used to generate a comparison CMD; the sources within the detection and
refernce circle are compared as part of the assessment process described in
Sec. 4.3. The plot on the lower left shows the smoothed surface density map of
the CMD-selected stars, with the same magenta and yellow circles and the green
ellipse as are shown in the upper left plot. The CMD in the upper right shows
all the sources in the field (black points), the CMD filter (blue solid line)
and the CMD-selected stars (red points). The CMDs for objects located within
the the detection circle and comparison (reference) circle appear on the lower
right. The detection pipeline found Pal 3 at 100% significance at a distance
range of 86$-$106 kpc, which brackets the actual distance of the cluster. It
was also found at a distance of $\sim$160 kpc because of a combination of
factors, including the shape of the CMD filter; see Sec. 5.1.2 for a full
discussion.

Pal 3 is located in the outer stellar halo of the Milky Way and is one of only
five Galactic GCs with galactocentric distances of $\sim$90 kpc or greater
(Harris, 1996, 2010 edition). It is fairly faint for a GC, with $M_{V}$ $=$
$-$5.7 (Harris, 1996, 2010 edition), has a modest mass (1.9 x 104 $M_{\odot}$;
Baumgardt & Hilker (2018)), and is markedly extended in size, with a half-
light radius ($r_{h}$) of $\sim$0.7′(Baumgardt & Hilker, 2018), or $\sim$19 pc
at the 91.9 kpc distance. This is several times larger than the median half-
light radii of the GCs in both the Milky Way and Andromeda, which is
$\sim$2$-$3 pc (van den Bergh, 2010). CMD fitting indicates that the cluster
belongs to the metal-poor subpopulation of the Milky Way GC system, with
[Fe/H] $=$ $-$1.7, and has an estimated age of $\sim$10 Gyr (Hilker, 2006).
The location and orbital properties of Pal 3 have prompted studies of its
origin, with some concluding that it may have been accreted into the outer
Galactic halo along with its parent dwarf satellite galaxy (Palma et al.,
2002), and others suggesting it could have been ejected from the Phoenix dwarf
irregular galaxy and captured by the Milky Way (Sharina et al., 2018).
However, stellar abundance studies seem to suggest instead that Pal 3 co-
evolved with the rest of the Galactic GC system (Koch et al., 2009).

It has long been understood that GCs like Pal 3 and the other members of the
Galactic GC system may contain gas and dust that originated in the outer
atmospheres of evolved stars and was subsequently lost to the interstellar
medium within the cluster (e.g., Roberts, 1959, 1960). Much of the gas is
expected to be removed from the GCs through a variety of mechanisms, including
winds from stars and stellar remnants, and dynamical pressure as the cluster
orbits the Galaxy (e.g., Frank & Gisler, 1976, Spergel, 1991). However some
gas may remain and be detectable as, for example, diffuse atomic hydrogen that
emits 21-cm radiation, or perhaps hotter gas that produces UV or X-ray
emission.

Many searches have been carried out over the past several decades using a
range of approaches to look for hydrogen gas in and around Milky Way GCs; in
general these studies have had limited success, often resulting in upper
limits or ambiguous detections (e.g., Heiles & Henry, 1966; Kerr & Knapp,
1972; Knapp et al., 1973; Erkes & Philip, 1975; Birkinshaw et al., 1983;
Faulkner et al., 1991; Freire et al., 2001; van Loon et al., 2006, and many
others). One study of particular relevance is van Loon et al. (2009), which
was carried out as part of the Galactic Arecibo L-band Feed Array (GALFA)
survey. van Loon et al. (2009) observed four GCs and set 3-$\sigma$ limits of
6$-$51 $M_{\odot}$ on the amount of gas in the four clusters, and also
identified a compact high-velocity H i cloud roughly 1 degree away from the
outer halo GC Pal 4. They conclude that the cloud may be physically associated
with Pal 4 or may simply be a chance superposition, since the two objects
differ by $\sim$40 km s-1 in velocity and the cloud has a relatively large
separation ($\sim$2 kpc) and a high H i mass (3 x 106 $M_{\odot}$) if it is
actually at the same distance as Pal 4 (109 kpc).

In Figure 3, we show the location of Pal 3 and the UCHVC AGC 501816 in our ODI
images. The ALFALFA H i ellipse is marked, as are the tidal and half-light
radii of the globular cluster. We analyzed the ODI imaging data of this field
in two ways: with the stars that make up Pal 3 included in the source lists,
and with the stars in Pal 3 removed. The first approach provides a useful
check of our methods for finding metal-poor, old stellar associations in our
images, since it allows us to see how effectively the pipeline detects the Pal
3 stars, and whether it finds them at the appropriate distance. The second
approach allows us to carry out the usual search for stellar populations in a
wide area around AGC 501816, as we had done with all of the other UCHVCs in
the sample.

We first carried out the detection steps described in Sec. 4 with the Pal 3
stars in place. We used the same CMD filter as before, but we changed the
lower limit of the distance range to 25 kpc (down from 250 kpc) in order to
include Pal 3’s much closer distance in our search process. Searching the
field with the Pal 3 stars in place yielded multiple detections at the
location of Pal 3 with high statistical significance (100%). The distances
corresponding to this local maximum in significance ranged from 86$-$106 kpc,
a distance range that encompasses the actual measured distance of Pal 3 (91.9
kpc). For the detections within this distance range, the stars selected by the
CMD filter populated the RGB, HB, Main Sequence Turnoff (MSTO) and the upper
Main Sequence. The stellar densities for these detections were more than 15
times the mean density across the field. Figure 4 shows a typical example of
the detection results with the Pal 3 stars present in the images; in this
example, the stellar overdensity is detected at a distance of $\sim$93 kpc.

The pipeline also detected the Pal 3 stars at high significance (100%) at a
distance of $\sim$160 kpc. A close examination of the results shows that this
is caused by a combination of factors: the distance range being sampled, the
shape of the lower portion of the CMD filter, and the detection limits of our
data. When we use the CMD filter to search for stellar populations at close
distances – i.e., between $\sim$25 kpc and $\sim$150 kpc – the main sequence
turnoff (MSTO) portion of the filter is included in the filtering process.
This portion of the filter is wide compared to the other CMD features and
allows stars with a broader range of magnitudes and colors to be selected. At
nearby distances, the position of the MSTO feature also coincides with regions
of the CMD where our photometric errors are larger (i.e., at $i$ magnitudes
between $\sim$22$-$24). The end result is likely that a higher level of
contamination exists in the sample of filter-selected stars. The strong
detection at 160 kpc occurred because that is the point at which the broad
MSTO feature in the filter falls just below the typical detection limits of
our data, and therefore only stars that coincide with the narrower RGB and HB
portions of the filter are selected. This lowers the mean surface density of
stars selected across the field and allows the bright stars in Pal 3 that
intersect other parts of the filter to yield a strong detection with a high
surface density relative to the mean density in the field. This type of false
detection should not occur in any of our other searches, because we normally
only sample the distance range 250 kpc to 2.5 Mpc, which means that the MSTO
feature is well below our detection limits. The bottom line is that the
results of the search carried out with the Pal 3 stars in place confirms that
our detection method can easily find a genuine old, metal-poor stellar
population present in the images and that it yields a strong detection at the
correct distance.

We next removed Pal 3 by excluding stars within 3 $r_{h}$ of the cluster
center in order to search for stars that might be associated with the H i
source a few arc minutes away. A few overdensities with high significance
(${}_{>}\atop{{}^{\sim}}$90%) were identified, so we carried out the final set
of detection steps – i.e., cross-matching the CMD-filtered source list with
Gaia and SDSS to remove likely contaminants and then repeating the detection
process on the more restricted catalog. The process yielded two possible
stellar overdensities, both of which are marked in the map of the field in
Figure 3. One detection, with a final significance of 87.76%, closely
coincides with the H i centroid and is found at a distance of $\sim$90 kpc.
The detection is made up of only six stars that lie in the MS turnoff portion
of the CMD filter, which is the broadest part of the filter and also (at the
relevant distance) falls well below the 50% detection limit of our images. A
second detection has a statistical significance of only 80.04% and is made up
of only seven stars, but is worth mentioning because it is located very close
to Pal 3, and is well inside the tidal radius of the cluster. Five of the
seven stars lie in the MSTO region of the CMD filter and the others coincide
with the HB and lower RGB. The distance yielded by the CMD filter for this
object is $\sim$80 kpc.

The locations of the two detected overdensities are marked in the diagram of
the field presented in Figure 3. The CMDs for the sources that make up the
overdensities are shown in the lower panels of that figure. For reference, a
single isochrone with the age and metallicity of Pal 3 is plotted on each CMD,
shifted to the distance at which the stellar overdensity is detected.

The question that arises at this point is what these two modestly significant
detections might represent. One clear possibility is that the detected stellar
overdensities are not genuine stellar populations associated with Pal 3 or AGC
501816 but are instead simply chance superpositions of stars that happen to
have magnitudes and colors that fall within the CMD filter at the given
distances. Another option is that we have detected stars that are somehow
associated with Pal 3, with the H i source, or with both objects. The latter
might be possible if, for example, the H i and stars originated in Pal 3 and
then were subsequently removed somehow from the cluster, or the gas was
stripped first and then stars formed within it later.

Evidence that argues against the idea that the gas could have been stripped
from Pal 3 is that if the H i were at the same distance as Pal 3 and the
detected overdensities ($\sim$80$-$90 kpc), the mass of AGC 501816 (based on
combining the ALFALFA flux $S_{21}$ with the distance using Equation 7 in
Adams et al., 2013) would be $\sim$2$-$3 x 103 $M_{\odot}$. This is a
substantial amount of H i gas to have originated in a globular cluster,
especially since it is found at a projected separation of $\sim$7′, or
$\sim$180 pc at a distance of 90 kpc. Most other studies have estimated gas
masses on the order of tens of $M_{\odot}$ or less for globular clusters.
Another issue is that the gas appears at a velocity that makes it difficult to
distinguish from Galactic neutral hydrogen. As mentioned, the H i source was
included in the UCHVC sample slated for optical follow-up because of its other
properties and its proximity to Pal 3, and despite the fact that its
$|v_{LSR}|$ value, 99 km s-1, is lower than the 120 km s-1 minimum threshold.
The ALFALFA detection grid in the vicinity of AGC 501816 and Pal 3 shows
strong, widespread Galactic H i emission primarily at substantially lower
velocities, but also shows a modest amount of weak, diffuse Galactic emission
in the velocity range $\sim$100$-$125 km s-1; thus we cannot be certain that
AGC 501816 is not gas associated with the Milky Way. We conclude that
additional observations of both the H i and the optical sources in the field
around Pal 3 and AGC 501816 are needed in order to definitively determine
whether or not there is any connection between these two objects.

Figure 5: The CMDs and spatial positions of the point sources detected in the
images of five UCHVCs classified as ”Possible Detections”. The UCHVC name
appears in the upper left corner of the CMD. Each CMD shows the location of
the filter (blue solid line), the distance modulus and distance of the
detected stellar overdensity, and the 50% completeness level (green dashed
line). In the CMDs, the gray points are the point sources that survive cross-
matching with Gaia and SDSS, sources within 3′ of the overdensity peak are
plotted in black, and sources that also fall within the CMD filter are over-
plotted in red. In the spatial position plots, the gray points are point
sources that survive cross-matching with Gaia and SDSS and red points are
sources that are selected by the CMD filter. Also marked are the overdensity
peak (magenta circle of 3′ radius), the H i ellipse (green), and the reference
circle (yellow).

####  5.1.3 Five Objects Classified as ”Possible Detections”

Five of the UCHVCs we analyzed in the sample of 26 yielded stellar overdensity
detections that we categorize as ”Possible Detections” because they are only
modestly convincing, but may warrant future follow-up observations. These five
detections survived all of the analysis steps, including the cross-matching
step with the Gaia and SDSS catalogs, and still met our criteria for
identifying possible dwarf galaxy candidates. However, they were only
significant in one of the two smoothing kernels (in all cases, the 2′ kernel)
and some of their other properties suggest they may not be genuine stellar
counterparts to the UCHVCs. Figure 5 shows the CMDs and spatial locations of
the sources in the fields of each of the five UCHVCs and Table 3 provides
information about the detected stellar overdensity, including the
significance, estimated distance, the sky coordinates, the log of the H i gas
mass (if the UCHVC were at the same distance as the detected overdensity), the
estimated $M_{V}$ and stellar mass of the stellar counterpart (the calculation
of these quantities is described in Sec. 5.2), and the estimated ratio of the
atomic gas mass to the stellar mass ($M_{HI}/M_{*}$). Details about each of
these five objects are provided below.

AGC 115710:— This field yielded the stellar overdensity that had both the
highest significance (99.55%) and the closest CMD filter distance (270 kpc).
The CMD filter includes several bright objects that might be suitable targets
for follow-up spectroscopy. On the other hand, the overdensity is relatively
far from the H i centroid (7.4′ or $\sim$580 pc in projection), and a few
background galaxies appear in the area around the overdensity.

AGC 122835.— This UCHVC field yielded an overdensity with one of the highest
levels of significance, 95.30%. The same overdensity was detected both before
and after the catalog cross-matching step, and even increased in significance
between the two iterations. The distance to the stellar overdensity (based on
the location of the CMD filter) is estimated at 870 kpc and the projected
separation from the H i centroid is 4.5′, which corresponds to a physical
separation of 1.1 kpc at the 870 kpc distance. The stars in the detected
overdensity do a fair job of filling in the RGB region of the CMD filter, from
the top of the RGB down to the faint limits of the images. On the other hand,
the CMD of the sources in the reference circle looks similar, both in terms of
the number of stars and their locations in the color-magnitude plane.

AGC 208315.— The stellar overdensity detected near this UCHVC is only modestly
significant (86.87%) and has an estimated distance of 630 kpc. It is located
approximately 4.9′ ($\sim$900 pc) in projection from the H i centroid but
nevertheless lies almost entirely within the H i ellipse (Fig. 5). A small
collection of stars is readily visible in the spatial plot and the region
within the detection circle is more populated than the region within the
reference circle. However, the stars that make up the overdensity do not fill
in the RGB region of the CMD filter and all of them fall below the 50%
completeness line in the CMD.

AGC 249283.— The stellar overdensity associated with this UCHVC is of interest
mainly because it lies directly on top of the H i cloud, with a projected
separation from the H i centroid of 0.8′. This translates to $\sim$100 pc at
the estimated distance given by the CMD filter, which is 440 kpc. The stars
that make up the overdensity lie within both the RGB and HB regions of the CMD
filter, and the detection circle is clearly more populated than the reference
circle. The major drawback of this detection is its lower significance, which
decreased from $\sim$88% to 83.13% after the catalog-matching step.

AGC 749140.— The overdensity detected in this field has a significance of
93.04% and a CMD filter distance of 490 kpc, and became the prevalent
detection after the cross-matching step. It is about 3.7′ (527 pc) from the H
i centroid. The stars within the overdensity do appear more clustered than the
stars in the surrounding image; however, none of the stars in the CMD filter
are bright, and the detection is due to the presence of objects in the lower
part of the RGB only, where the photometric uncertainties begin to increase.

####  5.1.4 Object Classified as the ”Best Detection”

The images of the UCHVC AGC268071 yielded the most convincing dwarf galaxy
candidate identified in the ODI imaging data; this object falls in the ”Best
Detection” category. The initial run of the detection pipeline on the ODI
images of AGC 268071 yielded possible detections of an optical counterpart at
a range of distances with statistical significance above 90%, so we carried
out the catalog cross-matching steps and then searched with the more
restricted catalog. The end result was a detection of a stellar association
with high statistical significance that is located 9.2′ from the H i centroid.
The same detection appears at 95.7% significance when we use a 2′ smoothing
kernel and at 97.3% significance when we use a 3′ kernel. The estimated
distance for this putative dwarf galaxy counterpart is 570 kpc, with a
possible range between 490 and 590 kpc.

The diagram showing the results (including the CMDs for the field, the spatial
locations of the stars relative to the H i, and the smoothed stellar
distribution) for the 3′ smoothing kernel is shown in Figure 6. The stellar
overdensity detected with the 3′ smoothing kernel includes 11 stars that are
noticeably clustered in the spatial distribution on the sky and that populate
the RGB portion of the CMD filter reasonably well. Furthermore, there are
appreciably more stars in the detection circle (11 stars) compared to a
reference circle (four stars) of the same size placed at a random position
within the field (lower right of Fig. 6). On the other hand, there are a few
background galaxies that appear in the ODI image in the general area of the
overdensity, so it is certainly possible that the detection is at least
partially due to the presence of a collection of unresolved background
galaxies that happen to have magnitudes and colors that fall within the CMD
filter. Also concerning is the fact that the stellar overdensity is located
relatively far away from the H i centroid; at a distance of 570 kpc, the 9.2′
angular separation between the stellar overdensity and the H i translates to a
projected physical separation of 1.5 kpc. It may be worth noting that AGC
268071 has the largest W50 value of any of the UCHVCs included in our sample
(Table 1) and in the UCHVCs catalog published in Adams et al. (2013). Follow-
up spectroscopy of some of the stars in the detected overdensity would be
useful for determining whether the sources are genuinely associated with the H
i source found by ALFALFA, or are simply the result of a random clustering of
foreground and background objects. This would be challenging because the
detection CMD in Figure 6 shows that the stars in the detected overdensity are
faint, with the brightest having an $i$ magnitude of 20.4 mag.

Figure 6:  CMDs and spatial positions of the point sources detected in the
images of the UCHVC AGC 268071. A stellar overdensity with a statistical
significance of 97.30% is detected $\sim$9 arcmin from the centroid of the H i
source. The upper left plot shows the locations of the stars detected in the
field (gray points), the stars selected by the CMD filter (red points), the H
i ellipse from ALFALFA (green ellipse), the 3′-radius detection circle
(magenta) and a reference circle of the same size placed at a random location
in the outskirts of the field. The smoothed surface density map of the CMD-
filtered stars is shown on the lower left, with the detection and reference
circles and H i ellipse marked as in the upper left plot. The CMD in the upper
right shows all the sources in the field (black points), the CMD filter (blue
solid line) and the CMD-selected stars (red points). The CMDs for objects
located with the the detection circle and comparison (reference) circle appear
on the lower right.

###  5.2 Estimating the Optical Properties of the Candidate Stellar
Populations

We can utilize the information provided by the ODI imaging data to calculate
rough estimates of the optical brightness of the potential stellar populations
we have detected. For completeness, we carry out this calculation for all six
of the sources in the current sample that we classify as either ”possible
detections” or the ”best detection”. We calculate two estimates of the
photometric properties of the detected stellar overdensities and combine those
with the estimated distance to produce a faint and bright limit for the
magnitude of the dwarf galaxy candidate that may be associated with the H i.

To calculate the faint estimate, we simply combine the flux from each of the
CMD-selected stars that make up the overdensity. This is a conservative
estimate of the brightness and assumes that these are the only stars
associated with the H i (i.e., that there is no underlying fainter stellar
population present).

To arrive at a bright limit for the magnitude of the possible stellar
population, we carry out aperture photometry of the images at the location of
the stellar overdensity. First we mask out bright objects that are not part of
the detected overdensity (e.g., extended background galaxies and stars that
are much brighter than the CMD-selected stars) as well as background pixels
that deviate appreciably from the median sky background level in that part of
the image. We replace the masked pixels with the median sky value and then
measure the flux in the region of the overdensity. We set our aperture size to
be based on the apparent angular size of the nearby dwarf galaxy Leo T (Irwin
et al., 2007) if it were located at the physical distance of our detection.
Therefore, the aperture ranges in size from roughly 1.4 to 4 arc minutes in
diameter, given the range of estimated distances of the detected stellar
overdensities (Table 3).

We also calculate a $g-i$ color that corresponds to the faint and bright
estimates, by subtracting the faint $i$ magnitude from the faint $g$
magnitude, and doing the same thing for the bright magnitudes. We use the
computed $g,i$ magnitudes and colors, the estimated distance, and the
photometric conversion relations in Jester et al. (2005) to calculate the
total absolute magnitude in the Johnson $V$ filter, $M_{V}$.

We can combine these magnitude and color estimates to calculate at least an
approximate stellar mass range for the candidate stellar populations we have
identified. We combine the relations in Bell et al. (2003) with our magnitudes
and colors to estimate the possible range of mass-to-light ratios for each
object and then apply those values to the faint and bright magnitude limits to
calculate the range of stellar masses. We also calculate the ratio of the H i
mass to the stellar mass for the faint and bright cases.

The optical magnitude and color estimates, and the associated H i mass to
stellar light ratios, are included in Table 3. The faint and bright estimates
for a given object can vary from each other by a few magnitudes, which means
the optical brightness of the putative stellar counterparts are highly
uncertain. Nevertheless, having at least a rough estimate allows us to compare
the properties of these objects to the properties of simulated and observed
dwarf galaxies in the Local Group (see Sections 6.2.2 and 6.2.3).

###  5.3 Revisiting the Stellar Overdensities Identified in the UCHVC Sample
from J19

Because of the improvements we had made to the stellar overdensity search
process, it seemed appropriate to reprocess and reassess the five potential
stellar overdensities that were detected in the J19 sample. We carried out a
full analysis of the imaging data for those five objects, beginning with
running source detection and photometry on the combined $g$ and $i$ image
pairs and proceeding all the way through the Gaia and SDSS catalog cross-
matching steps and the final assessment of the results. Information about the
five UCHVCs that were reanalyzed in this way is listed in Table 4; the table
columns are the same as those in the corresponding tables for the primary
sample of UCHVCs (Tables 1 and 2).

After processing and analyzing these objects with the updated procedures, we
found that four out of the five possible overdensities identified by J19 no
longer had convincing detections. Although there were a few cases where a
stellar overdensity with a significance of $\sim$80$-$90% or above was found
somewhere in the images, further scrutiny showed that in all cases the
overdensity was too far away from the H i source, and/or the CMD of the stars
within the detection circle looked too much like the CMD of the stars in the
reference circle. Accordingly, we now categorize these objects as non-
detections and are mark them as ”ND” in Table 4. Our overall conclusion is
that the changes to the detection pipeline described in Sec. 4.5 – especially
the reduction in foreground and background objects from our star lists made
possible by the catalog cross-matching – resulted in a more stringent process,
which in turn showed that four of the five most significant overdensities
presented in J19 are not likely to be genuine dwarf galaxies.

One UCHVC field, however, yielded a highly statistically significant dwarf
galaxy candidate that meets our criteria for a ”best detection”. The UCHVC is
AGC 249525, and the corresponding optical counterpart was first highlighted in
Janesh et al. (2017) and was one of the most convincing dwarf galaxy
candidates identified in the full sample of objects presented in J19. Using
our updated detection methods, we find a highly significant overdensity that
matches the characteristics of the original detection.

AGC 249525 has an H i mass of $\sim$107 $M_{\odot}$ (see Table 4), a column
density log($N_{HI}$) of 19.60 atoms cm-2 and an apparent HI major axis of
$\sim$9 arcmin. This source is one of the UCVHCs that was observed with the
WSRT and the VLA and has therefore been mapped at higher resolution (Adams et
al., 2016; Paine et al., 2020; Bralts-Kelly et al., 2020). Bralts-Kelly et al.
(2020) analyzed the VLA observations of AGC 249525 and concluded that they
show a velocity gradient that is indicative of rotation support, but that more
work is necessary to draw firmer conclusions about the H i kinematics of this
object.

The stellar overdensity detected with the earlier version of the pipeline and
presented in Janesh et al. (2017) and Janesh et al. (2019) had a statistical
significance of $\sim$98%. It was located well within the ALFALFA H i ellipse
and was directly coincident with the highest-density contour in the H i
synthesis maps from the WSRT presented in Adams et al. (2016). Janesh et al.
(2017) reported a distance of 1.60 Mpc, with an error (which they estimated by
determining the range of distances over which the significance stayed above
95%, for this object) of $\pm$0.45 Mpc. Their estimated stellar mass for the
optical counterpart was between $\sim$2 x 104 and 4 x 105 $M_{\odot}$.

Our updated analysis yields a possible stellar counterpart at the same
location as the one found by Janesh et al. (2017), but with an increased
statistical significance of 99.26%. The range of distances over which the
detected overdensity has a significance above 90% is 1.93$-$2.12 Mpc, and the
highest significance corresponds to a distance of 2.09 Mpc. This distance is
slightly larger than the J19 distance but the two distance ranges overlap.

The overdensity found by the new analysis is made up of 25 point sources that
are selected by the CMD filter and located within a region 3′ in radius. These
25 sources all fall within the upper portion of the RGB in the CMD filter (see
Figure 7). The number of sources in this detected overdensity is more than
twice the number of sources that make up the other overdensities in the sample
of UCHVCs presented in the current paper (see Table 3). The H i contours from
the WSRT observations presented in Adams et al. (2016) are over-plotted on the
spatial plots of the field that are shown in Figure 7; the position of the
stellar overdensity found with the updated pipeline remains coincident with
the H i contour with the highest $N_{\rm HI}$ value.

A distance of 2 Mpc is near the upper end of the distance range over which we
search for stellar counterparts, and at these large distances the upper
portion of the RGB coincides with the region of the CMD where the photometric
errors become larger, our detection completeness is lower, and it becomes more
difficult to distinguish true point sources from extended background objects
in our ground-based images. Therefore follow-up observations – deeper, higher-
resolution imaging and/or spectroscopy of the sources that make up the
detection – are needed in order to explore whether this overdensity is a
genuine stellar counterpart to the UCHVC or is simply a collection of
foreground and/or background objects that happen to coincide with the H i and
lie in the relevant portion of the CMD.

We carried out the steps described in Section 5.2 to estimate the optical
properties and the distance-dependent H i properties of the possible stellar
counterpart to AGC 249525; these are listed in Table 5. The $V$-band absolute
magnitude for the stellar counterpart is estimated to be in the range $-$5.99
to $-$6.87 mag. As one might expect (since the analysis was carried out on the
same images, and we find the overdensity at a similar distance), these
brightness estimates overlap the original brightness range estimated in J19
($M_{V}$ $=$ $-$4.5 to $-$7.1 mag).

Figure 7:  CMDs and spatial positions of the point sources detected in the
images of the UCHVC AGC 249525, which was included in the sample presented in
Janesh et al. (2019) and which we have reprocessed and analyzed with our
updated detection pipeline. A stellar overdensity with a statistical
significance of 99.26% is detected that coincides with the location of the H i
source (green contours). The upper left plot shows the locations of the stars
detected in the field (gray points), the stars selected by the CMD filter (red
points), the 3′-radius detection circle (magenta) and a reference circle of
the same size placed at a random location along the edges of the field. The
smoothed surface density map of the CMD-filtered stars appears on the lower
left, with the detection and reference circles marked in the same way as in
the upper left plot. The CMD in the upper right shows all the sources in the
field (black points), the CMD filter (blue solid line) and the CMD-selected
stars (red points). The CMDs for objects located with the the detection circle
and comparison (reference) circle appear on the lower right. The H i contours
from the Adams et al. (2016) WSRT study are shown with with green solid lines
on the upper left and lower left spatial plots; the contour levels are [9, 15,
20, 30, and 40] x 1018 atoms cm-2.

##  6 Summary and Discussion

###  6.1 Summary of Results

The UCHVCs investigated in this study were selected from the ALFALFA survey
because they had relatively compact angular sizes, were isolated from other H
i sources in both velocity space and location on the sky, had measured
velocities that suggested they were within the Local Volume and were not
likely to be Galactic H i, and did not have a clear optical counterpart in
existing catalogs and survey data. Our goal was to determine whether any of
the UCHVCs might actually host an as-yet undiscovered dwarf galaxy, perhaps
even one like the dwarf galaxy Leo P, which was first identified as an ALFALFA
UCHVC and then imaged in the optical with WIYN (Giovanelli et al., 2013; Rhode
et al., 2013). If the UCHVCs do indeed host a dwarf galaxy, they would be some
of the faintest, lowest-mass neutral-gas-bearing dwarf galaxies known in the
nearby Universe, with absolute V-band magnitudes in the range $\sim$$-$2 to
$-$8, stellar masses in the range $\sim$103 to 106 $M_{\odot}$, but with a gas
content in the range $\sim$105 to 106 $M_{\odot}$ (for an assumed distance of
1 Mpc).

In the current paper, we have presented results from the second and final
phase of our campaign to image a large sample of ALFALFA UCHVCs with WIYN. For
this phase of the campaign, we made several improvements to our procedures for
finding possible stellar counterparts to the H i clouds and we analyzed high-
quality $g-$ and $i-$band imaging of 26 UCHVCs to look for dwarf galaxy
candidates.

From the sample of 26 objects, we identified six UCHVCs with at least one
significant detection. We also carried out a detailed analysis of ODI imaging
of a UCHVC located a few arc minutes away from the Galactic outer halo GC Pal
3. Our best dwarf galaxy candidate is associated with the UCHVC AGC 268071;
the stellar overdensity we identify has a statistical significance of 97.30%,
a distance of $\sim$570 kpc, and is located roughly 9′ ($\sim$1.5 kpc) from
the centroid of the H i source.

We used our improved methods to reprocess and analyze the original WIYN
imaging of the five UCHVCs from the first phase of the project that had
possible stellar counterparts. From this set of five UCHVCs, we find one
stellar counterpart that qualifies as a convincing detection based on our new
procedures and criteria. The detected stellar overdensity appears at the same
location and a similar distance ($\sim$2 Mpc) as the stellar overdensity
originally identified in Janesh et al. (2017), but with an even higher
statistical significance (99.3%) when compared to repeated experiments that
randomly distribute the same number of stars around the field.

###  6.2 Putting Our Results in Context

####  6.2.1 Other Observational Searches for Dwarf Galaxy Counterparts to
Compact H i Clouds

In the time since the discovery of Leo P, we have had only limited success at
finding candidate dwarf galaxy counterparts to the UCHVCs selected from the
ALFALFA survey data, and we have certainly found no counterparts that are as
obviously apparent as Leo P. Leo P subtends roughly 90 arcsec on the sky, has
a total absolute $V$-band magnitude of $M_{V}$$\sim$ $-$9.3, and has an
underlying population of old ($\sim$12 Gyr) RGB stars, a young ($\leq$300 Myr)
population of massive early-type stars, a prominent HII region (Rhode et al.,
2013; McQuinn et al., 2015), and H-$\alpha$ emission that reveals the presence
of extended ($\sim$100 pc) ringlike structures (Evans et al., 2019) in the
galaxy’s interstellar medium. Given its properties, Leo P was readily visible
in our modest-length exposures (20-30-minute integrations in each of the
$B,V,R$ filters) taken with WIYN (Rhode et al., 2013).

Other studies have looked for counterparts to H i sources identified in
ALFALFA and other neutral hydrogen surveys and have also met with limited or
no success at finding Local Group dwarf galaxies, although they have
identified new dwarf galaxies within the Local Volume and/or at Virgo Cluster
distances. The survey that is most relevant to mention in this context is the
Galactic Arecibo L-band Feed Array HI (GALFA-HI) survey. GALFA-HI was aimed
primarily at studying the neutral hydrogen component of the Galactic
interstellar medium (Peek et al., 2011), but in the process identified 27
clouds that were classified as ”Galaxy Candidates” because of their kinematics
and separation from known galaxies or high-velocity cloud complexes (Saul et
al., 2012). Eleven of the ALFALFA UCHVCs from Adams et al. (2013) also appear
in the GALFA-HI catalog.

Several groups carried out optical follow-up observations and archival imaging
searches of objects chosen from the GALFA-HI and/or ALFALFA surveys.
Bellazzini et al. (2015a, b) used the Large Binocular Telescope to conduct the
SECCO Survey, which obtained deep optical images of 25 of the ALFALFA UCHVCs.
The survey resulted in the detection of a distant stellar counterpart to the
UCHVC AGC 226067 that was dubbed SECCO 1 (Bellazzini et al., 2015a).
Bellazzini et al. (2015b) suggested that the object was at least 3 Mpc away
and was most likely a star-poor dwarf galaxy in the Virgo Cluster. Adams et
al. (2015) subsequently published H i synthesis maps and optical imaging of
AGC 226067 and its putative optical counterpart. Tollerud et al. (2015) used
the WIYN pODI instrument to obtain follow-up imaging of 22 of the 27 GALFA-HI
Galaxy Candidates from the list published by Saul et al. (2012). From this
sample of 22 objects, Tollerud et al. (2015) identified two dwarf galaxies,
dubbed Pisces A and Pisces B, that were visible in SDSS images and appeared to
be nearby (at least within the Local Volume); they used spectroscopy with the
Palomar 5-m to confirm that the dwarfs were genuinely associated with the H i.
A subsequent HST imaging study by Tollerud et al. (2016) placed Pisces A and
Pisces B at distances of 5.6 Mpc and 8.9 Mpc, respectively (we note that these
objects were also cataloged by ALFALFA and designated AGC 103722 and AGC
114843; Haynes et al., 2018).

Sand et al. (2015) searched archival UV and optical data from a number of
different surveys and facilities (DSS, SDSS, Subaru SMOKA, CFHT Megacam,
GALEX, and Swift) to look for counterparts to a sample of UCHVCs selected from
GALFA-HI and ALFALFA. They looked for blue, diffuse emission similar to the
faint blue emission that was visible in the SDSS images of Leo P before its
discovery (Giovanelli et al., 2013; Rhode et al., 2013). They identified six
possible counterparts to their sample of UCHVCs and used spectroscopic follow-
up to confirm that five of the counterparts were genuinely associated with the
HI gas. All five of these counterparts are dwarf galaxies that lie outside the
Local Group. Two of the dwarfs were the above-mentioned galaxies Pisces A and
Pisces B from Tollerud et al. (2015). A third dwarf galaxy was SECCO 1, the
object identified by Bellazzini et al. (2015a); Sand et al. (2017) followed up
with an HST study that showed that this object is $\sim$17 Mpc away and is
likely a remnant dwarf galaxy produced through a ram pressure stripping event
that occurred within the M86 subgroup of the Virgo Cluster. The remaining two
dwarf galaxies found by Sand et al. (2015), GALFA Dw3 and GALFA Dw4, were
observed with HST to confirm their location outside the Local Group (at 7.6
Mpc and 3.1 Mpc, respectively) and note their isolated nature (no other
galaxies within 1.5 Mpc of either dwarf; Bennet et al., 2022).

Tollerud & Peek (2018) utilized the Exploring the Local Volume In Simulations
(ELVIS) suite of simulations (Garrison-Kimmel et al., 2014) to create mock H i
catalogs and explore how these compare to the results from GALFA-HI and the
associated optical searches for Local Group dwarf galaxies. The ELVIS
simulations were dark-matter-only simulations designed to model environments
similar to the Local Group, with two massive galaxy halos in the same
configuration (e.g., mass and position) as the dark matter halos that are
thought to host M31 and the Milky Way. Tollerud & Peek (2018) utilized a
simple empirical model to translate the dark-matter-only ELVIS simulations
into GALFA-HI observables and then examine what the H i survey would be
expected to find. Based on these mock catalogs, they found that they should
potentially discover tens of Local Group dwarf galaxies (depending on the
model), in contrast to the zero that were found in the actual observations.
They argued that this discrepancy could be explained if reionization inhibited
star formation at mass scales below the mass of galaxies like Leo T, which had
an estimated virial mass at the time of reionization of $\sim$108.5
$M_{\odot}$.

The results shown in Tollerud & Peek (2018) were further investigated by
DeFelippis et al. (2019), who combined the data in the GALFA-HI catalog with
archival optical imaging data from the Panoramic Survey Telescope and Rapid
Response System (Pan-STARRS; Chambers et al. (2016)). DeFelippis et al. (2019)
began by searching the GALFA-HI data for H i sources with the expected neutral
gas properties of nearby dwarf galaxies and identifying 690 objects. They then
applied a filtering technique to the Pan-STARRS archival data to look for
stellar populations with a range of ages and metallicities that might be
spatially associated with the H i. They also tested their algorithm by
searching for, and successfully finding, known dwarf galaxies (namely, Leo T
and Draco) at the appropriate distances. They found one potential dwarf galaxy
candidate near the Galactic plane, but no objects that resemble Leo T, and
argued that the results rule out the existence of Leo T-like objects within
the GALFA footprint (which covers $\sim$one-third of the sky) at distances
within 500 kpc. They conclude that their results support the Tollerud & Peek
(2018) assertion that reionization caused the gas to be removed from Local
Group objects with halo masses less than the $\sim$108.5 $M_{\odot}$
threshold, and thereby prevented those objects from developing into full-
fledged dwarf galaxies.

####  6.2.2 Comparison of Our Results to Model Simulations

A useful next step for understanding our results in a larger context is to
compare our most convincing dwarf galaxy candidates to the results from recent
state-of-the-art simulations that attempt to model the formation of dwarf
galaxies down to the lowest masses. Applebaum et al. (2021) (hereafter A21)
presented results from the DC Justice League suite of Milky Way-like zoom-in
simulations, which were designed to probe galaxy formation down to the regime
of ultra-faint dwarf galaxies (UFDs), i.e., down to absolute magnitudes of
$M_{V}$ fainter than $-$8 and stellar masses ${}_{<}\atop{{}^{\sim}}$105
$M_{\odot}$. The simulations were initially run to model large ($\sim$50 Mpc
across), dark-matter-only volumes and then small regions within those volumes
were re-simulated at higher resolution with full hydrodynamical treatment. The
smaller regions were specifically chosen to contain Milky Way analogs to
enable meaningful comparisons between the simulation results and the observed
properties of our local galaxy environment. Two Milky Way analogs, nicknamed
”Elena” and ”Sandra”, with masses of 7.5 x 1011 $M_{\odot}$ and and 2.4 x 1012
$M_{\odot}$ respectively, were selected for detailed simulations because they
were isolated and had masses consistent with current observational estimates
of the Milky Way’s virial mass. The DC Justice League simulations included
star and gas particles, feedback from supernovae and stellar winds, chemical
evolution, star formation, and supermassive black hole (SMBH) formation and
feedback. A21 cataloged the dwarf galaxies that were produced in each
simulation and determined the satellite galaxies’ stellar, gas, and virial
masses, luminosities, half-light radii, metallicities, galactocentric
distances, and star formation time scales. They were able to resolve UFDs with
absolute magnitudes as faint as MV $=$ $-$3 and found that they could match
the range of dynamical properties and size-luminosity scaling relations seen
in the known Local Group dwarf galaxies down to scales of $\sim$200 pc. They
also compared the properties of their simulated dwarf galaxies to the expected
detection limits that the Legacy Survey of Space and Time (LSST) data from the
Rubin Observatory will reach after 10 years. They concluded that virtually all
Local Group UFD galaxies will be detectable by LSST down to the luminosity
limit probed by their simulations.

Figure 8:  A recreation of Figure 5 from Applebaum et al. (2021), which shows
$V$ magnitude as a function of normalized distance for the dwarf galaxies in
the Sandra and Elena simulations (blue and orange points, respectively). The
dwarf galaxy candidates associated with the UCHVCs AGC 268071 and AGC 249525
are plotted on the figure (points with large error bars) for comparison. For
both simulations, satellites are denoted with filled squares, field galaxies
are shown as open squares, and backsplash galaxies are denoted with a star.
For the optical counterparts to AGC268071 and AGC 249525, we plot the average
of the ratios calculated from the bright and faint estimates of their
photometric properties and use both estimates to bracket our uncertainty. The
distance uncertainty for these two objects is determined by the distance over
which each object is detected in our data (see Table 3). We note that the
error bars representing the distance uncertainties for our dwarf galaxy
candidates are not symmetric in this plot and in Figures 9 and 10. This
reflects the fact that the significance of a given overdensity falls below the
90% or 80% threshold relatively quickly after the peak significance is
reached, as the tip of the RGB in the CMD filter is shifted downward to
fainter magnitudes and the handful of stars that make up a typical detected
overdensity become too bright to coincide with the filter.

Because the DC Justice League simulations were able to produce dwarf galaxies
with realistic properties, it should be instructive to compare the properties
of their dwarf galaxies with the properties of the two most convincing dwarf
galaxy candidates we identified in our search, namely the optical counterparts
we find for the UCHVCs AGC 268071 and AGC 249525. We have recreated two of the
figures from A21 that feature observable quantities that are also available
for our dwarf galaxy candidates. In Figure 8, which is a recreation of Figure
5 from A21, we show the absolute V-band magnitudes of the simulated dwarf
galaxies as a function of the distance to their parent galaxy normalized by
the parent galaxy’s virial radius. The figure makes it clear that the Sandra
host galaxy was associated with many more satellite galaxies than the Elena
host galaxy. A21 investigated this difference and attributed it to both the
higher mass of the Sandra halo and to the fact that an analog to the Large
Magellanic Cloud (which is associated with many of the ultra-faint dwarfs in
the vicinity of the Milky Way) appears in the Sandra simulation but is not
present in the Elena simulation. The simulated galaxies are also separated
into three categories: satellite galaxies that are within the virial radius of
their parent galaxy, field galaxies that are outside the virial radius of the
parent, and ”backsplash” galaxies that are field galaxies that have
experienced an infall that has brought them within their parent galaxy’s
virial radius.

Figure 9:  A recreation of Figure 13 from A21, which shows the H i-to-stellar
mass ratio as a function of normalized distance from the parent galaxy for
simulated dwarf satellite galaxies associated with Milky Way analogs. We have
added the dwarf galaxy candidates we found that are associated with the UCHVCs
AGC 268071 and AGC 249525. Blue points represent galaxies in the Sandra
simulation, while orange points represent galaxies in the Elena simulation.
Note that many of the galaxies in the simulations do not have H i gas and
therefore are not plotted. For AGC 268071 and AGC 249525 (points with large
error bars), we plot the average of the bright and faint magnitude estimates
and use both estimates to set the uncertainty. The distance uncertainty for
these two objects is determined by the distance range over which we detect
them (see Table 3).  Figure 10: A plot of dwarf galaxy $M_{V}$ values, drawn
from McConnachie (2012), as a function of distance from the parent galaxy for
galaxies in and near the Local Group. Milky Way satellite galaxies are plotted
in blue while M31 satellites are shown in orange. Dwarf galaxies that are not
satellites of either the Milky Way or M31 appear in pink, including Leo T
(pink filled star symbol). For the latter set of objects, we plot the distance
to the barycenter of the Local Group. Leo P is marked with a filled gray star
and plotted using the distance from McQuinn et al. (2015). We also show our
measurements for the dwarf galaxy candidate associated with AGC 268071 and AGC
249525 (black and red points with large error bars). For the two dwarf galaxy
candidates, we plot the average of the bright and faint absolute magnitude
estimates and use both estimates to set the uncertainty. The distance
uncertainty is defined by the distance range over which we detect the stellar
overdensity. We also note that the distances we are plotting for these two
objects are the heliocentric distances.

We have added our two dwarf galaxy candidates to Figure 8. To include them in
the plot, we calculated the distances of the two dwarf galaxy candidates in
terms of virial radius of the Milky Way (200 kpc; Dehnen et al., 2006). In the
figure, AGC 268071 appears to fit within the population of field galaxies,
which follow a weak trend of having fainter absolute magnitudes with
increasing distance from the parent galaxy. One of the simulated field
galaxies that is produced in the Sandra simulation has properties that overlap
those of AGC 268071, given the uncertainties on our estimated absolute
magnitude and distance for this dwarf galaxy candidate. We include AGC 249525
for completeness, although it lies in a region that is not populated with the
dwarf galaxies included in the simulation; it has an estimated absolute
magnitude that is similar to those of AGC 268071 and the other field dwarf
galaxies, but a distance that corresponds to more than 10 times the virial
radius of the Milky Way.

In Figure 9, we present a similar figure to Figure 13 in A21, and show the H
i-to-stellar mass ratios of the gas-bearing dwarf galaxies that appear in the
Sandra and Elena simulations. Only the galaxies that A21 determined to have
retained their gas (i.e., those that are not classified as ”quenched” in their
paper) are included in our version of the figure. The ratios are shown as a
function of the distance (again normalized by the virial radius) from the
parent galaxy. We also show the locations of the dwarf galaxy candidates we
detect in the images of the UCHVCs AGC 249525 and AGC 268071. Many of the
dwarf galaxies produced by the DC Justice League simulations do not have any
neutral gas and therefore are not plotted. The dwarf galaxy candidate
associated with AGC 268071 exists much farther out in distance from the parent
galaxy than the majority of the simulated dwarfs and lies on the high end of
the range of H i-to-stellar mass ratios. The dwarf galaxy candidate
corresponding to AGC 249525 is even more extreme, in terms of both gas-to-
stellar ratio and distance from the massive galaxy. Both dwarf candidates
appear generally consistent with a mild trend for higher gas-to-stellar ratios
with larger distances that is present in the simulated sample. The low stellar
content (and/or high H i-to-stellar mass ratio) of the dwarfs at large
distances may be at least partially explained by their isolated nature, as
they would presumably be less subject to tidal effects from their parent
galaxy or other nearby satellites and therefore perhaps less likely to either
be stripped of their gas and/or to have undergone starbursts triggered by
close interactions.

####  6.2.3 Comparison of Our Results to the Population of Local Group Dwarf
Galaxies

In addition to comparing to simulated data, it is worthwhile to compare the
dwarf galaxy candidates we detect to the actual dwarf galaxy population in and
around the Local Group. In Figure 10, we plot the absolute $V$-band magnitude
as a function of distance for dwarf galaxies within $\sim$3 Mpc. The values we
show are taken from the sample included in McConnachie (2012), who compiled
the measured properties (e.g., distances, magnitudes, metallicities,
structural and dynamical characteristics) of galaxies in the Local Group and
its environs. In the figure, Milky Way satellite galaxies are shown in blue,
M31 satellites are shown in orange, other nearby dwarf galaxies are shown in
pink, and our dwarf galaxy candidates are shown in black and red. The
positions of Leo T and Leo P are plotted with filled star symbols that are
pink and gray, respectively. Since the other nearby dwarf galaxies are not
bound to either the Milky Way or M31, we use their distance from the
barycenter of the Local Group as the distance value on the x-axis of the
figure, except that for Leo P we use the distance from McQuinn et al. (2015).
For the counterparts to AGC 268071 and AGC 249525, we simply use the estimated
heliocentric distance determined from our detection and analysis process.

Our estimate of the absolute $V$-band magnitude of the dwarf galaxy candidate
associated with AGC 268071 falls toward the faint end of the range of absolute
magnitudes for Milky Way satellites and at the very bottom end of the range of
luminosities of M31 satellites. The main difference between the AGC 268071
dwarf galaxy candidate and the known satellites of the Milky Way and M31 is
its large distance from the Milky Way (which we estimate at $\sim$570 kpc).
This distance places it substantially farther from the Milky Way or from M31
than any of the other dwarf satellite galaxies that appear on the plot. This
suggests that it may not be a satellite of the Milky Way, but may instead be
simply a Local Group dwarf galaxy (if it is a genuine galaxy at all); in that
case, the most relevant comparison sample for this object would be the other
nearby dwarf galaxies that are not satellites (pink points with distances
extending to $\sim$3 Mpc on Figure 10). Compared to this sample of objects,
the dwarf galaxy candidate associated with AGC 268071 is roughly 2 magnitudes
less luminous, and is also closer in distance than all but a few of the dwarf
galaxies in the figure. Overall it seems that the detection for AGC 268071
makes a reasonable dwarf galaxy candidate in the sense that it at least lies
within a similar parameter space as the other dwarf galaxies in and around the
Local Group, although on the low-luminosity edge of the distribution.

The estimated $M_{V}$ magnitude and distance for the dwarf galaxy candidate
associated with AGC 249525 puts this galaxy $\sim$2.5 magnitudes below that
same sample of nearby dwarf galaxies. This object would thus be a gas-rich
ultra-faint dwarf galaxy that is substantially fainter than the other
currently-known galaxies at this distance ($\sim$2 Mpc). As noted in Janesh et
al. (2017), a potential close neighbor to AGC 249525 in terms of distance and
sky position is UGC 9128, which is 10 degrees away on the sky and has a
heliocentric velocity $cz$ of 152 km s-1(McConnachie, 2012), compared to 48 km
s-1 for AGC 249525.

###  6.3 Next Steps and Future Observing Capabilities

In our analysis of the ALFALFA UCHVCs, we found six possible dwarf galaxy
candidates. However, we again note that – even with our deep optical multi-
band imaging data, with typical seeing $\sim$1 arcsec or better – there was no
clearly visible diffuse optical emission, nor any obvious (to the eye) stellar
association, at the locations of the dwarf galaxy candidates detected via the
CMD filtering process. Without such additional evidence, it is appropriate to
question whether the objects we found are indeed genuine dwarf galaxies. The
statistically-significant stellar overdensities we have detected may actually
be made up of some combination of compact background galaxies and Galactic
foreground stars (that are below the detection limits for Gaia EDR3 and/or
SDSS and therefore too faint to be eliminated in the cross-matching step)
which happen to coincide with the rough location of the H i source. To confirm
that the stars in the detected overdensities are truly associated with the
corresponding UCHVC, we would need spectroscopic follow-up observations that
show that the stars’ velocities are similar to that of the H i. The brightest
stars that make up the detections in our best two candidates, AGC 268071 and
AGC 249525, have $i$ magnitudes of 20.4 and 23.1 mag, respectively. Deriving
accurate radial velocities of stars of these magnitude, even with time on an
$\sim$8$-$10-meter telescope, is challenging; for example, our initial attempt
to obtain a spectrum of a star in the AGC 268071 dwarf galaxy candidate using
the LRS2 instrument (Chonis et al., 2016) on the 10-meter Hobby-Eberly
Telescope555Based on observations obtained with the Hobby-Eberly Telescope
(HET), which is a joint project of the University of Texas at Austin, the
Pennsylvania State University, Ludwig-Maximillians-Universitaet Muenchen, and
Georg-August Universitaet Goettingen. The HET is named in honor of its
principal benefactors, William P. Hobby and Robert E. Eberly. yielded a
spectrum with insufficient signal-to-noise to measure a radial velocity. The
rest of the stars that make up the overdensities for our two best candidates
are even fainter (see Figures 6 and 7). Thus, follow-up spectroscopy is
extremely challenging with current ground-based optical spectroscopy
capabilities.

More generally, further progress toward completing the census of low-mass
galaxies in the Local Group and its environs will be made with the advent of
the next generation of large-scale systematic surveys like the Rubin
Observatory’s Legacy Survey of Space and Time (LSST; Ivezić et al. 2019). LSST
will uniformly survey the Southern sky in six filters ($ugrizy$) over the
course of 10 years; the median 5-$\sigma$ point source detection limit should
be $r=24$ mag in the single-visit images and $r$ $\sim$27 mag in the 10-year
co-added images, with median effective seeing of 1.0 arcsec. With this wide-
field, multi-color, high-quality imaging data, LSST should be able to resolve
RGB stars out to $\sim$10 Mpc within its survey area. The survey should be
sensitive to galaxies as faint as the candidate dwarf galaxy we detected in
the AGC 268071 field (i.e., $M_{V}$ $\sim$ $-$6) in galaxy groups out to
$\sim$3$-$4 Mpc (LSST Science Collaboration et al., 2009; Simon, 2019). LSST
should also be able to probe even fainter galaxies with resolved stellar
photometry out to 1$-$2 Mpc and should be complete for galaxies with stellar
luminosities ${}_{>}\atop{{}^{\sim}}$2 x 103 $L_{\odot}$ that are within
$\sim$1 Mpc of the Milky Way (Bullock & Boylan-Kolchin, 2017). Comparing the
catalogs of resolved stars and newly-discovered dwarf galaxies that will
likely be detected by LSST to H i observations from sensitive radio facilities
like the Five Hundred Meter Spherical Telescope (FAST), the MeerKAT telescope,
the ASKAP telescope, and the planned Square Kilometre Array Observatory (SKAO)
may help reveal the nature of objects like the UCHVCs or the Galaxy Candidates
identified in GALFA-HI data by Saul et al. (2012) and whether these are
genuine gas-rich ultra-faint dwarf galaxies or simply isolated gas clouds with
no associated stellar population.

The authors wish to thank the anonymous referee for providing helpful
suggestions that improved the manuscript. We also thank the staff of the WIYN
Observatory and Kitt Peak National Observatory for their assistance during our
pODI and ODI observing runs. We acknowledge the work of the ALFALFA
collaboration in observing, processing, and extracting H i sources. We are
grateful to the staff members at WIYN, NOIRLab, and Indiana University
Pervasive Technology Institute for designing and implementing the ODI
Pipeline, Portal, and Archive (ODI-PPA) and assisting us with the pODI and ODI
image reduction. K.L.R., N.J.S., and W.F.J. were supported by NSF grant
AST-1615483. We thank the Indiana University (IU) College of Arts and Sciences
for funding IU’s share of the WIYN telescope. This research has made use of
the NASA/IPAC Extragalactic Database (NED), which is funded by the National
Aeronautics and Space Administration and operated by the California Institute
of Technology.

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Table 1: Sample of UCHVCs Analyzed for This Work

Name | R.A. | Dec. | $S_{21}$ | $cz$ | $w_{50}$ | $\bar{a}$ | log$\bar{N}_{HI}$ | log$M_{HI}$ | Obs. Date | Detection
---|---|---|---|---|---|---|---|---|---|---
|  |  | [Jy km s-1] | [km s-1] | [km s-1] | [′] | [atoms cm-2] | [M⊙] |  |
AGC102994 | 00h54m31.6s | +29d24m02s | 0.58 $\pm$ 0.03 | -290 | 23 $\pm$ 2 | 5.70 | 18.90 | 5.14 | 2016B, 2017B | ND
AGC115710 | 01h26m21.5s | +05d23m08s | 4.81 $\pm$ 0.04 | 59 | 21 $\pm$ 1 | 8.65 | 19.45 | 6.05 | 2016B | PD
AGC116548 | 01h06m00.7s | +12d26m57s | 0.96 $\pm$ 0.03 | 84 | 11 $\pm$ 1 | 6.15 | 19.05 | 5.35 | 2015B | ND
AGC122834 | 02h03m48.8s | +29d13m13s | 3.58 $\pm$ 0.06 | 49 | 16 $\pm$ 5 | 10.70 | 19.14 | 5.93 | 2016B | ND
AGC122835 | 02h05m35.5s | +29d13m56s | 1.23 $\pm$ 0.04 | 29 | 23 $\pm$ 5 | 5.60 | 19.24 | 5.46 | 2016B | PD
AGC174764 | 07h56m14.8s | +25d09m00s | 0.64 $\pm$ 0.04 | 174 | 21 $\pm$ 1 | 7.00 | 18.76 | 5.18 | 2016B | ND
AGC198683 | 09h32m08.0s | +23d37m52s | 0.88 $\pm$ 0.04 | 178 | 19 $\pm$ 1 | 11.10 | 18.50 | 5.32 | 2016B | ND
AGC208315 | 10h27m01.1s | +08d47m08s | 4.96 $\pm$ 0.07 | 148 | 20 $\pm$ 2 | 12.60 | 19.14 | 6.07 | 2017A | PD
AGC208524 | 10h47m02.5s | +01d46m31s | 0.67 $\pm$ 0.03 | 179 | 14 $\pm$ 1 | 4.85 | 19.10 | 5.20 | 2016A, 2017A | ND
AGC208752 | 10h23m09.0s | +20d40m59s | 1.79 $\pm$ 0.03 | -60 | 16 $\pm$ 1 | 7.60 | 19.13 | 5.63 | 2019A | ND
AGC219214 | 11h09m29.8s | +05d26m01s | 0.56 $\pm$ 0.03 | 142 | 20 $\pm$ 5 | 5.55 | 18.90 | 5.12 | 2019A | ND
AGC227977 | 12h09m20.0s | +04d23m30s | 0.57 $\pm$ 0.05 | -143 | 18 $\pm$ 4 | 10.00 | 18.40 | 5.13 | 2014A | ND
AGC227988 | 12h46m22.9s | +04d48m42s | 0.49 $\pm$ 0.05 | 320 | 48 $\pm$ 10 | 4.70 | 18.99 | 5.06 | 2019A | ND
AGC233763 | 13h12m42.3s | +13d30m46s | 0.78 $\pm$ 0.04 | 129 | 31 $\pm$ 6 | 5.30 | 19.09 | 5.26 | 2017A | ND
AGC233831 | 13h22m41.6s | +11d52m31s | 0.63 $\pm$ 0.02 | 124 | 16 $\pm$ 1 | 4.50 | 19.14 | 5.17 | 2018A | ND
AGC249283 | 14h23m57.7s | +05d23m40s | 1.11 $\pm$ 0.07 | 252 | 32 $\pm$ 9 | 13.50 | 18.43 | 5.42 | 2019A | PD
AGC257994 | 15h53m54.0s | +14d41m48s | 2.04 $\pm$ 0.04 | 146 | 23 $\pm$ 3 | 9.65 | 18.98 | 5.68 | 2021A | ND
AGC268067 | 16h05m29.4s | +16d09m12s | 1.94 $\pm$ 0.05 | 158 | 35 $\pm$ 7 | 8.40 | 19.08 | 5.66 | 2018A | ND
AGC268071 | 16h12m36.8s | +14d12m26s | 2.67 $\pm$ 0.08 | 109 | 62 $\pm$ 15 | 9.75 | 19.09 | 5.80 | 2018A | BD
AGC268213 | 16h22m35.7s | +05d08m48s | 2.90 $\pm$ 0.06 | -139 | 18 $\pm$ 4 | 10.85 | 19.03 | 5.83 | 2017A | ND
AGC333604 | 23h11m23.2s | +27d56m45s | 1.74 $\pm$ 0.03 | 66 | 19 $\pm$ 2 | 7.65 | 19.12 | 5.61 | 2018B | ND
AGC333651 | 23h55m21.4s | +25d17m26s | 0.87 $\pm$ 0.03 | 45 | 14 $\pm$ 4 | 7.85 | 18.79 | 5.31 | 2015B | ND
AGC334257 | 23h02m11.3s | +16d00m48s | 0.68 $\pm$ 0.04 | -452 | 22 $\pm$ 11 | 7.95 | 18.68 | 5.20 | 2018B | ND
AGC335755 | 23h14m16.4s | +03d23m07s | 7.95 $\pm$ 0.10 | 42 | 23 $\pm$ 1 | 13.15 | 19.31 | 6.27 | 2018B | ND
AGC501816 | 10h05m19.4s | -00d02m26s | 1.70 $\pm$ 0.05 | 107 | 20 $\pm$ 1 | 7.15 | 19.17 | 5.60 | 2019A | ND
AGC749140 | 00h51m16.4s | +15d11m11s | 1.33 $\pm$ 0.03 | 52 | 17 $\pm$ 4 | 5.95 | 19.22 | 5.50 | 2018B | PD

Note. — Following Haynes et al. (2018), the adopted error on the recessional
velocities listed in column 5 is half of the w50 error. The H i mass values in
column 9 are calculated at an assumed distance of 1 Mpc. In the ”Detection
Status” column, objects for which no detection in the field meets our lower-
threshold criteria are listed as ND; this includes the UCHVC AGC 501816, which
is located near the globular cluster Pal 3, and which is discussed in Sec.
5.1.2. Objects categorized as having a possible detection are listed as PD,
and objects with our best (most convincing) detections are listed as BD; see
Sec. 4.3 for more information about the categories.

Table 2: UCHVCs With Insufficient Data to Include in the Sample

Name | R.A. | Dec. | $S_{21}$ | $cz$ | $w_{50}$ | $\bar{a}$ | log$\bar{N}_{HI}$ | log$M_{HI}$ | Obs. Date
---|---|---|---|---|---|---|---|---|---
|  |  | [Jy km s-1] | [km s-1] | [kms-1] | [′] | [atoms cm-2] | [M⊙] |
AGC219663 | 11h34m29.7s | +20d12m49s | 0.75 $\pm$ 0.03 | 74 | 17$\pm$1 | 7.35 | 18.79 | 5.25 | 2013A
AGC229327 | 12h32m31.6s | +17d57m21s | 1.02 $\pm$ 0.05 | 251 | 25$\pm$1 | 12.95 | 18.43 | 5.38 | 2013A
AGC232765 | 13h23m09.4s | +15d11m17s | 1.37 $\pm$ 0.03 | 105 | 23$\pm$3 | 5.15 | 19.36 | 5.51 | 2016A
AGC249326 | 14h31m58.8s | +06d35m20s | 0.70 $\pm$ 0.04 | 136 | 38$\pm$11 | 5.65 | 18.98 | 5.22 | N/A
AGC249393 | 14h10m58.1s | +24d12m04s | 1.01 $\pm$ 0.06 | -156 | 36$\pm$1 | 11.90 | 18.50 | 5.38 | 2013A
AGC249441 | 14h07m00.8s | +00d13m23s | 0.74 $\pm$ 0.04 | -131 | 16$\pm$2 | 8.40 | 18.66 | 5.24 | g: 2019A
AGC249565 | 14h35m57.6s | +17d10m04s | 1.76 $\pm$ 0.04 | 30 | 18$\pm$1 | 7.65 | 19.12 | 5.62 | 2013A
AGC257956 | 15h55m07.5s | +14d29m29s | 1.54 $\pm$ 0.04 | 144 | 25$\pm$6 | 6.70 | 19.18 | 5.56 | i: 2017A
AGC258241 | 15h08m24.4s | +11d24m22s | 0.98 $\pm$ 0.07 | 163 | 19$\pm$2 | 12.70 | 18.43 | 5.36 | N/A

Note. — The adopted error on recessional velocities listed in column 5 is half
of the w50 error according to Haynes et al. (2018). The H i mass values in
column 9 are calculated at an assumed distance of 1 Mpc.

=60mm

Table 3: Stellar Overdensities Found in the UCHVC Images Analyzed for This Work Name | Significance | DistanceaaThe distance listed here is the CMD filter distance associated with the most significant detection of a given overdensity. The numbers within the parentheses represent the range of distances for which a given detection remains above some significance threshold (90% for overdensities that have a peak significance $\geq$90%, and 80% for overdensities with a peak significance between 80$-$90%.) | R.A. | Dec. | log MHI | MV | MV | $g-i$ | $g-i$ | log M∗ | log M∗ | MHI/M∗ | MHI/M∗
---|---|---|---|---|---|---|---|---|---|---|---|---|---
| (# stars) | (Range) |  |  |  | (faint) | (bright) | (faint) | (bright) | (faint) | (bright) | (faint) | (bright)
| [%] | [Mpc] |  |  | [M⊙] | [mag] | [mag] | [mag] | [mag] | [M⊙] | [M⊙] |  |
Possible Detections |  |
AGC115710 | 99.55(9) | 0.27 (0.25-0.31) | 01:25:53.7 | +05:25:51.3 | 4.93${}^{+0.11}_{-0.08}$ | -2.18${}^{+0.19}_{-0.27}$ | -7.01${}^{+0.19}_{-0.27}$ | 0.92$\pm$0.02 | 1.19$\pm$0.01 | 3.22${}^{+0.11}_{-0.08}$ | 5.35${}^{+0.11}_{-0.08}$ | 50.82$\pm 0.12$ | 0.38$\pm 0.01$
AGC122835 | 95.30(10) | 0.87 (0.86-0.92) | 02:05:26.0 | +29:17:57.1 | 5.34${}^{+0.05}_{-0.01}$ | -4.47${}^{+0.03}_{-0.12}$ | -7.67${}^{+0.03}_{-0.12}$ | 1.57$\pm$0.03 | 1.49$\pm$0.07 | 4.62${}^{+0.05}_{-0.01}$ | 5.50${}^{+0.05}_{-0.01}$ | 5.17$\pm 0.01$ | 0.69$\pm 0.01$
AGC208315 | 86.87(12) | 0.63 (0.59-0.64) | 10:27:08.8 | +08:51:41.2 | 5.67${}^{+0.02}_{-0.01}$ | -3.63${}^{+0.14}_{-0.04}$ | -6.97${}^{+0.14}_{-0.04}$ | 0.26$\pm$0.03 | 0.63$\pm$0.08 | 3.30${}^{+0.02}_{-0.05}$ | 4.78${}^{+0.02}_{-0.05}$ | 232.23$\pm 0.54$ | 7.74$\pm 0.02$
AGC249283 | 83.13(9) | 0.44 (0.42-0.50) | 14:23:58.7 | +05:22:56.6 | 4.72${}^{+0.10}_{-0.04}$ | -3.52${}^{+0.11}_{-0.24}$ | -8.15${}^{+0.11}_{-0.24}$ | 1.41$\pm$0.02 | 1.59$\pm$0.05 | 4.13${}^{+0.10}_{-0.04}$ | 5.74${}^{+0.10}_{-0.04}$ | 3.91$\pm 0.01$ | 0.09$\pm 0.01$
AGC749140 | 93.04(9) | 0.49 (0.43-0.49) | 00:51:20.7 | +15:14:46.3 | 4.88${}^{+0.01}_{-0.11}$ | -3.18${}^{+0.27}_{-0.02}$ | -6.92${}^{+0.27}_{-0.02}$ | 0.84$\pm$0.03 | 0.94$\pm$0.05 | 3.56${}^{+0.01}_{-0.11}$ | 5.13${}^{+0.01}_{-0.11}$ | 20.64$\pm 0.10$ | 0.56$\pm 0.01$
Best Detection |  |
AGC268071 | 97.30(11) | 0.57(0.49-0.59) | 16:12:55.6 | +14:20:45.3 | 5.31${}^{+0.04}_{-0.13}$ | -4.29${}^{+0.32}_{-0.10}$ | -6.99${}^{+0.32}_{-0.10}$ | 1.12$\pm$0.01 | 1.03$\pm$0.01 | 4.21${}^{+0.04}_{-0.13}$ | 5.31${}^{+0.04}_{-0.13}$ | 12.48$\pm 0.03$ | 1.21$\pm 0.01$
Table 4: UCHVCs from J19 That Were Reanalyzed for This Study

Name | R.A. | Dec. | $S_{21}$ | $cz$ | $w_{50}$ | $\bar{a}$ | log$\bar{N}_{HI}$ | log$M_{HI}$aaThe distance listed here is the CMD filter distance associated with the most significant detection of the overdensity. The numbers within the parentheses represent the range of distances for which the detection remains above 90% significance.  | Obs. Date | Detection
---|---|---|---|---|---|---|---|---|---|---
|  |  | [Jy km s-1] | [km s-1] | [km s-1] | [′] | [atoms cm-2] | [M⊙] |  |
AGC198606 | 09h30m05.5s | +16d39m03s | 6.73 $\pm$ 0.67 | 53 | 21$\pm$1 | 9.00 | 19.60 | 6.20 | 2013A | ND
AGC215417 | 11h40m08.1s | +15d06m44s | 0.70 $\pm$ 0.07 | 216 | 17$\pm$4 | 9.49 | 18.50 | 5.20 | 2013A | ND
AGC219656 | 11h51m24.3s | +20d32m20s | 0.85 $\pm$ 0.08 | 192 | 21$\pm$1 | 8.00 | 18.80 | 5.30 | 2013A | ND
AGC249525 | 14h17m50.1s | +17d32m52s | 6.73 $\pm$ 0.67 | 48 | 24$\pm$7 | 9.00 | 19.60 | 6.20 | 2013A | BD
AGC268069 | 16h05m32.6s | +14d59m20s | 1.14 $\pm$ 0.11 | 132 | 29$\pm$4 | 7.07 | 19.00 | 5.40 | 2013A | ND

Note. — Following Haynes et al. (2018), the adopted error on the recessional
velocities listed in column 5 is half of the w50 error. The H i mass values in
column 9 are calculated at an assumed distance of 1 Mpc. In the ”Detection
Status” column, objects for which no detection in the field meets our lower-
threshold criteria are listed as ND, and objects with our best (most
convincing) detections are listed as BD; see Sec. 4.3 for more information
about the categories.

=60mm

Table 5: Stellar Overdensities Found in the Reanalyzed UCHVC Images from J19 Name | Significance | DistanceaaThe distance listed here is the CMD filter distance associated with the most significant detection of the overdensity. The numbers within the parentheses represent the range of distances for which the detection remains above 90% significance.  | R.A. | Dec. | log MHI | MV | MV | $g-i$ | $g-i$ | log M∗ | log M∗ | MHI/M∗ | MHI/M∗
---|---|---|---|---|---|---|---|---|---|---|---|---|---
| (# stars) | (Range) |  |  |  | (faint) | (bright) | (faint) | (bright) | (faint) | (bright) | (faint) | (bright)
| [%] | [Mpc] |  |  | [M⊙] | [mag] | [mag] | [mag] | [mag] | [M⊙] | [M⊙] |  |