# A Deep Census of Outlying Star Formation in the M101 Group

Ray Garner, III  Department of Astronomy, Case Western Reserve University,
10900 Euclid Avenue, Cleveland, OH 44106, USA  J. Christopher Mihos
Department of Astronomy, Case Western Reserve University, 10900 Euclid Avenue,
Cleveland, OH 44106, USA  Paul Harding  Department of Astronomy, Case Western
Reserve University, 10900 Euclid Avenue, Cleveland, OH 44106, USA  Aaron E.
Watkins  Astrophysics Research Institute, Liverpool John Moores University,
Liverpool, L3 5RF, UK  Centre for Astrophysics Research, School of Physics,
Astronomy & Mathematics, University of Hertfordshire, Hatfield, AL10 9A, UK

###### Abstract

We present deep, narrowband imaging of the nearby spiral galaxy M101 and its
group environment to search for star-forming dwarf galaxies and outlying H II
regions. Using the Burrell Schmidt telescope, we target the brightest emission
lines of star-forming regions, H$\alpha$, H$\beta$, and [O III], to detect
potential outlying star-forming regions. Our survey covers
$\sim$$6\text{\,}{\mathrm{deg}}^{2}$ around M101, and we detect objects in
emission down to an H$\alpha$ flux level of
$5.7\text{\times}{10}^{-17}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$
(equivalent to a limiting SFR of
$1.7\text{\times}{10}^{-6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$
at the distance of M101). After careful removal of background contaminants and
foreground M stars, we detect 19 objects in emission in all three bands, and 8
objects in emission in H$\alpha$ and [O III]. We compare the structural and
photometric properties of the detected sources to Local Group dwarf galaxies
and star-forming galaxies in the 11HUGS and SINGG surveys. We find no large
population of outlying H II regions or undiscovered star-forming dwarfs in the
M101 Group, as most sources ($93\text{\,}\mathrm{\char 37\relax}$) are
consistent with being M101 outer disk H II regions. Only two sources were
associated with other galaxies: a faint star-forming satellite of the
background galaxy NGC 5486, and a faint outlying H II region near the M101
companion NGC 5474. We also find no narrowband emission associated with
recently discovered ultradiffuse galaxies and starless H I clouds near M101.
The lack of any hidden population of low luminosity star-forming dwarfs around
M101 suggests a rather shallow faint end slope (as flat as $\alpha\sim-1.0$)
for the star-forming luminosity function in the M101 Group. We discuss our
results in the context of tidally-triggered star formation models and the
interaction history of the M101 Group.

††facilities: CWRU:Schmidt††software: Astropy (Astropy Collaboration et al.,
2013, 2018), Matplotlib (Hunter, 2007), NumPy (Harris et al., 2020)

##  1 Introduction

Understanding star formation and the variety of locales in which it takes
place is key to understanding galaxy formation and evolution. Star formation
is most easily seen and studied in the inner luminous regions of galaxies
(e.g. Martin & Kennicutt, 2001), but star formation in low-density
environments is less well understood, whether that be in low luminosity dwarf
galaxies (Kennicutt et al., 2008; Lee et al., 2007, 2009) or outlying H II
regions (Rudolph et al., 1996; Ferguson et al., 1998; Lelièvre & Roy, 2000;
Ryan-Weber et al., 2004; Werk et al., 2010). In recent years, progress has
been made to observationally explore these environments in ultraviolet and
optical light.

UV investigations of the outer regions of galaxies has been primarily aided by
the _Galaxy Evolution Explorer_ (_GALEX_) satellite. _GALEX_ revealed that
over $\sim$$30\text{\,}\mathrm{\char 37\relax}$ of spiral galaxies posses UV-
extensions of their optical disks, possibly indicating that low-density star
formation is not rare (Thilker et al., 2007). Moreover, extended-UV (XUV)
emission and outyling H II regions beyond the optical radius of the disk are
associated with previous or ongoing galaxy interactions (Thilker et al., 2007;
Werk et al., 2010).

In optical light, outlying isolated H II regions have been detected via
narrowband imaging targeting specific emission lines, primarily H$\alpha$,
where they appear as emission-line point sources. These regions have been
found in environments ranging from galaxy clusters and compact groups to the
halos of galaxies (e.g. Sakai et al., 2002; Mendes de Oliveira et al., 2004;
Ryan-Weber et al., 2004; Walter et al., 2006; Boquien et al., 2007; Werk et
al., 2010; Kellar et al., 2012; Keel et al., 2012). These H II regions
indicate recent formation of OB stars outside of the normal star-forming
environment of the inner disk. And unlike the XUV emission of spiral galaxies,
these outlying H II regions can appear well outside twice the canonical
$R_{25}$ size of the nearest galaxy (Ryan-Weber et al., 2004; Werk et al.,
2010). This gives us insight into extreme modes of star formation and may
contribute to intragroup and intracluster light (see Vílchez-Gómez 1999 for a
review).

Possible sites for this low-density star formation are dwarf galaxies. Dwarf
galaxies are very common around massive galaxies and in group environments.
Star forming dwarf galaxies (SFDGs) are one of the most common types of
galaxies in the local universe (e.g. de Lapparent, 2003). They are
characterized by low stellar mass, low chemical abundance, high gas content,
and high dark matter content (see Gallagher & Hunter 1984 for a review). They
tend to lie in low density environments (Weisz et al., 2011a, b) and are
typically of low surface brightness. The low surface density of cold gas in
SFDGs is below that at which star formation is truncated (the so-called
Kennicutt-Schmidt law, e.g. Schmidt 1959; Kennicutt 1989, 1998), but star
formation is not completely halted (Hunter et al., 1998). These galaxies are
noted for their rather chaotic spatial distributions of their star forming
regions, typically being asymmetrical and clumped on large scales (e.g. Hodge,
1975). The extremely low surface brightnesses of many of these systems makes
detection of the continuum stellar emission difficult, and SFDGs would likely
appear as clumps of bright emission-line sources in H$\alpha$ surveys.

Additionally, types of dwarf galaxies are being found with very low masses
(ultrafaint dwarfs, UFDs; Simon 2019) and very low surface brightness
(ultradiffuse dwarfs, UDGs; Sandage & Binggeli 1984; van Dokkum et al. 2015).
Both types of galaxies represent the extreme faint end of the galaxy
luminosity function, typically having luminosities fainter than $M_{V}=-7.7$
(Simon, 2019). These dwarf galaxies tend to have ancient ages often consistent
with star formation ending by reionization at $z\sim 6$ (Brown et al., 2014).
If there are star-forming analogues to these galaxy types in the local
universe, they may be detectable through very deep and wide-field narrowband
imaging (van der Hulst et al., 1993; McGaugh & Bothun, 1994; Schombert et al.,
2011; Cannon et al., 2011, 2018).

In an effort to explore these areas of low-density star formation, we have
used Case Western Reserve University’s Burrell Schmidt 24/36-inch telescope to
perform the first deep, wide-field, multiline, narrowband observations of the
nearby spiral galaxy M101 (NGC 5457, $D=$6.9\text{\,}\mathrm{Mpc}$$; see
Matheson et al. 2012 and references therein) and its group environment. We
image in three different emission lines that are characteristic of star
formation, H$\alpha$, H$\beta$, and [O III], which allows us to more cleanly
reject contaminants without the need for expensive follow-up spectroscopy.
This strategy, combined with our large survey area
($\sim$$6\text{\,}{\mathrm{deg}}^{2}$) and our deep photometric limit of
$5.7\text{\times}{10}^{-17}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$
(reaching SFR limits of
$1.7\text{\times}{10}^{-6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$)
gives us a good census of extended and outlying star formation and faint star-
forming dwarf galaxies over large areas in nearby groups.

M101 was chosen for this survey because its nearby distance enables its
properties to be studied in detail (Mihos et al., 2012, 2013, 2018; Watkins et
al., 2017). M101 is also currently interacting with its satellite population,
likely its massive satellite NGC 5474 as evidenced by its asymmetric disk
(Beale & Davies, 1969; Rownd et al., 1994; Waller et al., 1997). Given that
interacting systems often display extended or outlying star-forming regions,
the M101 Group could be an excellent case study to explore the conditions
under which extended intragroup star formation is triggered. Additionally,
M101 has been found to harbor ultradiffuse galaxies (Merritt et al., 2014,
2016; Karachentsev et al., 2015; Danieli et al., 2017; Carlsten et al., 2019)
and constraining the star-forming properties of these objects will aid in
understanding star formation in low-density environments.

##  2 Narrowband Imaging

The narrowband imaging used here was taken over the course of three seasons
using Case Western Reserve University’s 24/36-inch Burrell Schmidt telescope,
located at Kitt Peak in Arizona. Our narrowband imaging and data reduction
techniques are described in detail in Watkins et al. (2017) and summarized
briefly here.

The Burrell Schmidt images a $$$\times$$$ field of view onto a single
4096$\times$4096 back-illuminated CCD, yielding a pixel scale of $None$
${\mathrm{pixel}}^{-1}$. For each emission line studied (H$\alpha$, H$\beta$,
and [O III]), we image in two narrowband filters (see Table 1) — one centered
on the emission line and another shifted
$\approx$$150\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ off the emission line
for continuum subtraction. Given the width of our filters (necessitated by the
fast $f/3.5$ beam of the Schmidt; Nassau 1945), our H$\alpha$-on filter covers
both H$\alpha$ and the adjoining [N II]$\lambda\lambda$6648,6583 lines, while
the [O III]-on filter covers both lines of the [O
III]$\lambda\lambda$4959,5007 doublet. In each filter, we image M101 using
55–71 images of $1200\text{\,}\mathrm{s}$ exposure time each, with each
pointing dithered randomly by up to $None$. All data was taken under dark,
photometric conditions, with the H$\alpha$ imaging taken in Spring 2014 (and
described in detail in Watkins et al. 2017), the H$\beta$ imaging in Spring
2018, and the [O III] imaging in Spring 2019. Flat fielding was done using a
combination of twilight flats and offset night sky flats (see Watkins et al.,
2017), and deep $1200\text{\,}\mathrm{s}$ observations of Regulus and Arcturus
were used to model and correct for scattered light from bright stars in the
field (see Slater et al., 2009). We subtract sky from each image using a
simple plane fit to the background sky levels across the image (typically
$\sim$$100\text{\,}\mathrm{ADU}$). After correcting for sky and scattered
light, as well as for fringing from OH sky lines in the H$\alpha$ on-band
filter, each set of dithered images was then median-combined to yield a final
on- and off-band master image of the field, representing a total exposure time
of 18–24 hours in each filter.

We flux calibrate the narrowband imaging using three methods. First, we
calibrate using observations of spectrophotometric standard stars (Massey et
al., 1988) taken throughout the course of each night, solving for extinction
coefficients and nightly zeropoints that are applied to each image. Second, we
self-calibrate each image using the 100–150 stars in the field around M101
that have well-measured photometry from Sloan Digital Sky Survey (SDSS)
imaging, applying a color-dependent offset between SDSS broadband filters and
our narrowband filters synthesized using the Pickles (1998) Stellar Spectral
Flux Library. Third, we self-calibrate each image using the the $\sim$100 SDSS
spectroscopic point sources in the field, using the SDSS spectroscopy to
synthesize calibrated AB magnitudes in our narrowband filters for zeropointing
each image. These three independent techniques yielded zeropoints which agreed
with one another to within $\pm$$5\text{\,}\mathrm{\char 37\relax}$, which we
take as the uncertainty in our absolute photometric calibration.

Table 1: Narrowband Imaging Datasets

Filter | $\lambda_{0}$ | $\Delta\lambda$ | Exposure Time | ZP (flux) | ZP (AB)
---|---|---|---|---|---
H$\alpha$-on | 6590 Å | 101 Å | $71\times$1200\text{\,}\mathrm{s}$$ | $5.61\text{\times}{10}^{-18}$ | 26.63
H$\alpha$-off | 6726 Å | 104 Å | $71\times$1200\text{\,}\mathrm{s}$$ | $5.50\text{\times}{10}^{-18}$ | 26.64
H$\beta$-on | 4875 Å | 82 Å | $59\times$1200\text{\,}\mathrm{s}$$ | $7.65\text{\times}{10}^{-18}$ | 26.73
H$\beta$-off | 4757 Å | 81 Å | $55\times$1200\text{\,}\mathrm{s}$$ | $7.91\text{\times}{10}^{-18}$ | 26.74
[O III]-on | 5008 Å | 102 Å | $67\times$1200\text{\,}\mathrm{s}$$ | $7.58\text{\times}{10}^{-18}$ | 26.91
[O III]-off | 5114 Å | 101 Å | $66\times$1200\text{\,}\mathrm{s}$$ | $7.37\text{\times}{10}^{-18}$ | 26.89

Note. — ZP (flux) converts 1 ADU to $\rm erg\ s^{-1}\ cm^{-2}$ in the master
images, while ZP (AB) converts to AB magnitudes.

##  3 Methods

###  3.1 Source Detection

Our final imaging dataset from the Burrell Schmidt consists of the narrowband
imaging described in Section 2, along with deep broadband imaging in
Washington $M$, similar to Johnson $V$, and a modified (bluer) Johnson $B$
filter from Mihos et al. (2013). We start by detecting sources on the
H$\alpha$ on-band image, using astropy’s PhotUtil package, specifically the
segmentation module (Bradley et al., 2019). This program detects sources as
objects that have a minimum number of connected pixels that are each greater
than the background threshold value. In our case, the threshold value above
which pixels would be marked as a detection was $3\sigma$ above the background
level on the H$\alpha$ on-band image after a two pixel Gaussian smoothing.
This sigma-clipping is conceptually similar to SExtractor’s $\kappa\sigma$
clipping (Bertin & Arnouts, 1996).

In order to avoid detecting random noise spikes or star-forming objects well
inside known bright galaxies, we masked several regions in the H$\alpha$ on-
band image. First, we masked a $750\text{\,}\mathrm{pixel}$ ($\sim$$None$)
border around the image where the background noise becomes dominant. We also
masked stars in the Tycho-2 Catalog (Høg et al., 2000) brighter than
$B_{T}=12.5$. Circular masks were applied to many of the galaxies in our
survey area corresponding to twice the $R_{25}$ isophotal radius, taken from
the RC3 catalog (de Vaucouleurs et al., 1991). This mask size follows from
what has been done previously by other authors in defining the boundary beyond
which lies outlying H II regions (Werk et al., 2010). The only exception to
this was M101 itself; Mihos et al. (2013) showed that the $R_{25}$ reported in
the RC3 significantly overestimates M101’s $\mu_{B}=25$ isophotal radius by as
much as a factor of two. We utilized the areal-weighted $R_{25}=$$$ for M101
reported in Mihos et al. (2013).

Having masked these regions, we then create a two-dimensional background
object to calculate the background sky level and its uncertainty. The sky
level was estimated in boxes of $100\times 100$ pixels with filter sizes of
$10\times 10$ pixels. Then, as mentioned before, we detected sources that were
$3\sigma$ above the background, resulting in $32\,439$ sources. We used the
default parameters of 32 multi-thresholding levels and a contrast of $0.001$
to deblend close or overlapping sources. The segmentation module uses a
combination of multi-thresholding and watershed segmentation to separate
overlapping sources, which, given the parameters above, results in a
segmentation map with a total of $35\,308$ sources. For context, SExtractor
utilizes only a multi-thresholding technique to deblend sources (Bertin &
Arnouts, 1996).

Using this segmentation map to define each source, we calculated photometric
and structural quantities for each source in each of the narrowband images as
well as in the broadband imaging of Mihos et al. (2013). Many of these were
default calculations for segmentation, including positions and fluxes. We also
calculated photometric errors, signal-to-noise, and AB magnitudes for each
object in each filter. In what follows, the magnitude in each filter will be
written as $m_{\lambda_{0}}$, where $\lambda_{0}$ is the central wavelength of
the filter. Additionally, we calculate flux differences, $\Delta
f=f_{\text{on}}-f_{\text{off}}$, and emission line equivalent widths (EWs) in
each pair of filters. We again note that our H$\alpha$ on-band filter bandpass
includes the [N II]$\lambda$,$\lambda$6549,6583 doublet, thus measuring
$\text{H}\alpha+[\text{{N~{}II}}]$. We have corrected these fluxes by adopting
the emission line ratio of [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et
al., 2000).

Given that the net flux in a filter pair can be either positive or negative,
we express the net flux in each band using asinh magnitudes (Lupton et al.,
1999):

$m(\Delta f)=\text{ZP}-2.5\log(b)-a\operatorname{arcsinh}\left(\frac{\Delta
f}{2b}\right)$

Here, $a\equiv 2.5\log e=1.08574$ is Pogson’s ratio (Pogson, 1856), $\Delta f$
is the flux difference of the source in a particular filter, and the zero
point, ZP, was chosen so that objects with $\Delta f=0$ had zero magnitude.
This system has the benefit that it can be calculated for any flux value,
negative or positive, and behaves smoothly as the flux drops through zero. The
softening parameter, $b$, was chosen to be the flux of an object with
$\text{S/N}=1$. The softening parameter and zero point were calculated per
filter set; specific values can be seen in Table 2. In this magnitude system,
objects with positive net flux (“in emission”) will have numerically negative
magnitudes, while objects with negative net flux (“in absorption”) will have
positive magnitudes. Going forward, we will refer to these “net flux”
magnitudes as $m(\Delta f)$ in each spectral line.

Table 2: Asinh Magnitude Parameters Filter Set | $b$ | ZP
---|---|---
| [${10}^{-17}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$] |
H$\alpha$ | $3.8$ | -$41.05$
H$\beta$ | $5.4$ | -$40.67$
[O III] | $5.3$ | -$40.69$

We also cross-matched our source list with those objects in the SDSS DR16
(Ahumada et al., 2020) and with _Gaia_ Early Data Release 3 (EDR3; Gaia
Collaboration et al. 2020). The _Gaia_ data will aid in rejecting interlopers
in our dataset, while the SDSS data provides additional photometry and
structural information for our final sample of objects.

At this point, of those $35\,308$ sources detected, $5442$
($\sim$$15\text{\,}\mathrm{\char 37\relax}$) were detected “in emission” (i.e.
have positive net flux at the $3\sigma$ level) in H$\alpha$ and form the
starting sample for our emission line source catalog. Only $450$
($\sim$$1\text{\,}\mathrm{\char 37\relax}$) were detected “in emission” in all
of the three narrowband filters. However, as we will show next, many of these
detections were not true narrowband emission features.

###  3.2 Investigating Potential Sources

At this point in our detection routine, we have a source catalog of any
astronomical source that has excess emission in H$\alpha$. This could include
objects such as H II regions or dwarf galaxies in the nearby universe,
background emission line objects, or objects such as M stars with molecular
absorption bands in our filters.

This is illustrated in Figure 1 where we plot our “net flux” magnitudes and
EWs against the signal-to-noise of each source in each filter set. Since we
are attempting to detect H II regions, we select for only those sources with
net positive H$\alpha$ flux differences, i.e. negative “net flux” magnitudes.
Most of our sources have very low signal-to-noise, so we use only those
sources detected with $\text{S/N}>3$. These plots are conceptually similar to
those in Figure 2 of Kellar et al. (2012) used to search for compact emission
line sources, except that in place of their “ratio” quantity, we use a true
signal-to-noise, and in place of their magnitude difference, we use our “net
flux” magnitudes.

Figure 1: Top row: The H$\alpha$, H$\beta$, and [O III] asinh net flux
magnitude as a function of signal-to-noise for each source observed in
emission in H$\alpha$ (negative $m(\Delta f)$). Bottom row: The narrowband
equivalent width (EW) as a function of signal-to-noise in each filter set. In
both rows, the vertical dashed line indicates $\text{S}/\text{N}=3$; sources
at lower S/N are considered undetected in our analysis. Black points show all
sources, orange crosses are sources detected in emission in all three lines,
and green crosses are sources detected in two lines (H$\alpha$ and [O III]).
See text for details.

The most notable aspect of Figure 1 is that both the H$\beta$ and [O III]
filters have sources in absorption despite selecting for objects that are only
in emission in H$\alpha$. We investigate this further using spectral synthesis
of various objects – stars, nearby galaxies, and high redshift objects –
through our filters to assess their behavior in our sample selection criteria.

To test how other stars would behave in our filters, we synthesized spectra
from the Stellar Spectral Flux Library (Pickles, 1998) through our filters.
The Stellar Spectral Flux Library was chosen because it offers a wide
distribution of stars of various spectral types, luminosity classes, and
metallicity while also covering a large wavelength range with uniform $R\sim
500$ spectral resolution. For nearby galaxies, we adopt the SDSS DR5 spectral
template for different galaxy types, synthesizing their EWs at zero redshift,
such that the emission lines fall within our filters. The synthesized EWs in
each filter are shown as a function of $B-V$ color in Figure 2.

Figure 2: Synthesized equivalent width as a function of color for stars in the
Pickles (1998) Stellar Spectral Flux Library (black points) and $z=0$ SDSS
galaxy spectral templates (blue triangles). The horizontal dashed line
indicates the equivalent width selection cut for our emission line sample in
each filter pair.

It is clear from Figure 2 that stars bluer than $B-V\sim 1$ will broadly mimic
H$\alpha$ emission, H$\beta$ absorption, and [O III] emission. Meanwhile,
stars redder than $B-V\sim 1$ will broadly mimic H$\alpha$ emission, H$\beta$
emission, and [O III] absorption. However, these are not true narrowband
absorption or emission features, but rather are a result of how our narrowband
filters sample features in the stellar continuum. Because continuum starlight
produces these low level “pseudo-emission” features, in our search for true
star-forming signatures, we ignore any object with EW values lower than
$\text{EW}(\text{H}\alpha)=$8\text{\,}\mathrm{\SIUnitSymbolAngstrom}$$,
$\text{EW}(\text{H}\beta)=$2\text{\,}\mathrm{\SIUnitSymbolAngstrom}$$, and
$\text{EW}(\text{[{O~{}III}]})=$5\text{\,}\mathrm{\SIUnitSymbolAngstrom}$$ as
shown by the dotted lines in Figure 2. Given the relative weaker strength of
the H$\beta$ line compared to H$\alpha$, we make a separate distinction
between sources with all three filters in emission, satisfying all the EW cuts
above, and sources with H$\alpha$ and [O III] filters in emission, satisfying
only those two EW cuts above. These two groups will be called the three-line
sample and two-line sample, respectively. Making those cuts reduces our source
catalog from $35\,308$ sources to $147$ sources with $\text{S/N}>3$ (95 in the
three-line sample and 52 in the two-line sample).

At zero redshift, galaxy SEDs show the expected trend between line emission
and color: bluer late-type galaxies are stronger in EW than the redder early-
type galaxies due to young stellar populations in the former. The EW cuts made
above will only cut out the reddest galaxies with little to no H$\alpha$ or [O
III] emission. Therefore these cuts will not negatively impact our search for
star-forming galaxies even at relatively low EW.

Finally, we investigate contamination of our sample due to high redshift
objects in the M101 field. In the H$\alpha$ survey of Kellar et al. (2012),
$37\text{\,}\mathrm{\char 37\relax}$ of their detected sources were higher
redshift objects where the [O III] line was redshifted into their H$\alpha$
filter. Watkins et al. (2017) detected a handful of background galaxies and
quasars in their H$\alpha$ sample, so the possibility of detecting high
redshift objects in any one filter is strong. While our use of three
narrowband filters significantly reduces the chance of background
contaminants, there are still regions of redshift space where bluer emission
lines can redshift into our filters.

Figure 3: Regions of redshift space where bright nebular emission lines
manifest emission or absorption signatures as they redshift through our
narrowband filters. Note that the legend indicates the narrowband filter pair,
not the redshifted emission line.

Figure 3 shows where common bright emission lines in star-forming galaxies can
shift through our filters over the range of redshift $0\leq z\leq 0.5$. At
$z<0.01$, all three filters will appear in emission. If we select for only
those redshift ranges that satisfy the EW cuts above, i.e. objects detected in
the three- or two-line samples, then emission line sources in the narrow
redshift window $0.314<z<0.332$ can also potentially contaminate our samples.
Here, the H$\alpha$ emission line has redshifted out of our filters entirely,
while the [O III]$\lambda\lambda 4959{,}5007$ doublet has redshifted into the
H$\alpha$ filters. Similarly, the H$\beta$ line has redshifted into the [O
III] filters and the [O II]$\lambda\lambda 3727{,}3729$ lines have redshifted
into the H$\beta$ filters. Due to the small size of the Burrell Schmidt, star-
forming galaxies at higher redshift are unlikely to be detected at all, but
bright high redshift AGN may still produce some contamination of the sample.
However, the rarity of such objects makes them unlikely to be present in large
numbers in our sample.

###  3.3 Removing M Stars

While the equivalent width cuts significantly reduce contamination due to
Milky Way stars, M stars continue to pose a particular challenge. The reason
for this can be seen in Figure 4, which shows the spectrum of a M5V star from
the Stellar Spectral Flux Library (Pickles, 1998). Overplotted are the filter
transmission curves of our narrowband filters. The H$\alpha$ off-band filter
lies in a TiO molecular absorption trough, producing net emission in the
measured H$\alpha$ flux. Other features in the complex stellar continuum mimic
emission in our H$\beta$ and [O III] filters as well. The abundance of
Galactic M stars at the faint end of the stellar mass function suggests many
of our “emission line” detections will be M star contaminants needing to be
removed.

Figure 4: The spectrum of a M5V star from Pickles (1998) (solid black line)
overplotted with our narrowband filter transmission curves (colored lines),
showing how features in the stellar continuum of M stars produce pseudo-
emission signatures in our narrowband filters.

One possible method to removing M stars from our detections is to use some
definition of compactness to distinguish between stellar point sources and
extended H II regions. Indeed, such a measure of compactness, the Gini index
(e.g. Lotz et al., 2004), is part of the standard calculations made by the
segmentation routine. However, given the FWHM of the Schmidt imaging ($None$)
and the assumed distance to M101, we would only be able to resolve objects
larger than $75\text{\,}\mathrm{pc}$. The Strömgren sphere radius of an H II
region powered by an O9 star is $\sim$$30\text{\,}\mathrm{pc}$, illustrating
that some H II regions would be unresolved in our imaging. Therefore, using
compactness to distinguish between stars and H II regions would bias us
against the detection of smaller H II regions in our survey.

Instead, to remove any remaining Milky Way stars from our emission line
sample, we cross-match our sources with the _Gaia_ EDR3 catalog. Since star-
forming objects in the M101 Group should have no detectable parallax or proper
motion, we cut any cross-matched source with a $3\sigma$ detection of parallax
or proper motion. The results of these cuts are shown in Figure 5. Nearly all
of the sources with $B-V>1$ are rejected by the _Gaia_ cuts on parallax and
proper motion, indicating that they are Galactic M stars. There are a few red
objects that do not appear in the _Gaia_ catalog at all. Of the brighter ones
at $V\sim 19$, the object at $B-V\sim 1.5$ is cataloged by SDSS as a galaxy
with a photometric redshift of $z=$0.188\pm 0.034$$, while redder one at
$V\sim 19$ and $B-V\sim 1.7$ is a point source in the SDSS imaging. Of the
three fainter objects at $V\sim 22-23$, the bluest one at $B-V=1.1$ is
cataloged as a galaxy in the SDSS imaging, with a photometric redshift of
$z=$0.388\pm 0.137$$. Given the uncertainty in the photometric redshift, this
object may be an example of a background contaminant leaking into our sample
through the redshift window shown in Figure 3. The other two faint red sources
undetected by _Gaia_ appear as point sources in the SDSS imaging. Based on
this combined analysis of _Gaia_ and SDSS properties, we reject all sources
redder than $B-V=1$, including the five mentioned here that do not appear in
the _Gaia_ catalog.

Figure 5: Broadband color-magnitude diagram for sources in both the three- and
two-line samples. The red points are those objects rejected by the _Gaia_
parallax and proper motion cuts. Characteristic photometric errorbars are
shown, which include both uncertainties in the source photometry and in the
overall photometric zeropoints.

###  3.4 Additional Cuts

Aside from contamination of the sample on the red end by Galactic M stars,
Figure 5 shows a handful of very blue objects at $B-V<0$. Cross-matching these
sources against the SDSS imaging catalog reveals they are background sources
in the M101 field. The two bluer objects are QSOs at redshift $z=1.34$ and
$z=0.76$, where bright emission lines from [O II] or Mg have shifted into our
filters. The reddest of the three is a background galaxy resolved in the SDSS
imaging, but lacks any spectroscopic or photometric redshift.

The high redshift objects mentioned above illustrate the need for a final cut
to remove them. For objects in emission in the H$\beta$ filters, we utilize
the $\text{H}\alpha/\text{H}\beta$ ratio as an additional interloper cut. For
unobscured ionized gas, the Balmer decrement should be
$\text{H}\alpha/\text{H}\beta=2.86$ (Osterbrock, 1989) with higher ratios
indicating higher extinction levels. Therefore objects in our sample which
show much lower Balmer decrements are likely background sources where other
emission lines have redshifted into the filters. Similarly, objects with
anomalously high Balmer decrements are likely also contaminants, as star-
forming galaxies in the local universe rarely get above a decrement of 8
(Domínguez et al., 2013). We therefore keep only objects with
$1<\text{H}\alpha/\text{H}\beta<8$ for the three-line sample. This cut removes
four objects, including one of the QSOs noted above.

There is no similar calculation we can make for those sources in the two-line
sample as they have no detected H$\beta$ emission. Additionally, depending on
the specific redshift and combination of lines moving through our filters, it
is possible that a background source could mimic a realistic Balmer decrement.
To combat this, we cross-matched with SDSS, removing those sources for which
SDSS has a spectroscopic or photometric redshift. This process removes 20
objects at higher redshift (8 QSOs and 12 galaxies), but also allows us to
reject $z\sim 0$ objects located just beyond M101. Two such examples of $z\sim
0$ contaminants were detected: the blue compact dwarf galaxy SBS 1407+540
($z=0.0068$, $d=$29.1\text{\,}\mathrm{Mpc}$$; Stepanian 2005) and the
Magellanic-type irregular galaxy CGCG 272-015 ($z=0.0071$,
$d=$30.4\text{\,}\mathrm{Mpc}$$; Ann et al. 2015). After all of the cuts, we
examined each object by eye to confirm the nature of the source; during this
process one object was discarded due to it being an obvious blend, pairing a
foreground M star with a background galaxy.

Summarizing, we have split our source catalog into two groups: the three-line
sample consists of those sources in H$\alpha$, H$\beta$, and [O III], while
the two-line sample is comprised of sources showing emission in H$\alpha$ and
[O III] only. Both samples have been cleaned of stellar contamination using a
combination of EW cuts and _Gaia_ parallaxes or proper motions. Background
contamination was removed by cross-matching with redshift estimates from SDSS
imaging and spectroscopy. Finally, for the three-line sample, an additional
cut was made on the observed Balmer decrement to reject non-physical values.
Our final samples consist of 19 objects in the three-line sample and 8 objects
in the two-line sample.

##  4 Analysis

In this section, we investigate and describe the properties of the sources in
the three-line and two-line samples. In total, the three-line sample contains
19 sources while the two-line sample contains 8 sources. Tables 3 and 4 list
the narrowband properties of the sources in the three-line and two-line
samples, respectively. The tables include a unique identifier for each source,
the right ascension and declination (epoch J2000), whether that source is
unresolved (U) or extended (E) in our images, the flux difference for each
filter pair, the equivalent width (EW) for each filter pair, and the
$\text{H}\alpha/\text{H}\beta$ and [O III]/H$\alpha$ ratios.

Figure 6: Broadband color-magnitude diagram for sources in the final three-
line (orange points) and two-line samples (green points). Circles show sources
unresolved in SDSS imaging, while triangles show extended sources.

Tables 5 and 6 list the broadband properties of the sources in the three-line
and two-line samples, respectively. Included are the source identifier, the
$V$-band magnitude, the $B-V$ color, the respective SDSS $ugriz$ photometry if
available, and the _GALEX_ far ultraviolet (FUV) and near ultraviolet (NUV)
magnitudes if available. The color-magnitude diagram of the sources is shown
in Figure 6.

Figure 7: (a) The broadband $B$ image with M101 and its companions, NGC 5474
and NGC 5477 labeled, as well as background galaxies also in the survey area.
(b) The continuum-subtracted H$\alpha$ image. The red circles represent the
masks used in our source detection algorithm. The orange and green squares
represent the objects in our three-line and two-line samples, respectively.
Both images measure $2.4\times 2.4$ $\mathrm{deg}$. North is up and east is to
the left. Figure 8: A zoomed-in view of the detected regions surrounding M101.
Left: the continuum-subtracted H$\alpha$ image. Right: the broadband $B$
image, stretched to show the faint structure of the outer disk. The orange and
green circles with radii of $None$ are centered on the objects in our three-
line and two-line samples, respectively. The red circle represents the size of
the mask used in our source detection algorithm. Both images measure $24\times
24$ $\mathrm{arcmin}$. North is up and east is to the left. Figure 9: A
zoomed-in view of the detected regions near NGC 5474 (top row) and NGC 5486
(bottom row). Left: the continuum-subtracted H$\alpha$ image. Right: the
broadband $B$ image, stretched to show the faint structure of the outer disks.
Colored circles are the same as in Figure 8. The images of NGC 5474 measure
$9\times 9$ $\mathrm{arcmin}$ and the images of NGC 5486 measure $3.6\times
3.6$ $\mathrm{arcmin}$. North is up and east is to the left.

Briefly, we give a sense of how deep we have attained in fluxes and the range
of equivalent widths investigated in each group. In the three-line sample, the
faintest source has $V=22.9$. The corresponding H$\alpha$ net flux is
$8\text{\times}{10}^{-16}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$,
which, at the $6.9\text{\,}\mathrm{Mpc}$ distance to M101 and using the
SFR-$L(\text{H}\alpha)$ calibration of Kennicutt & Evans (2012), would
correspond to a SFR of
$2.4\text{\times}{10}^{-5}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$.
The three-line sample has a range of H$\alpha$ equivalent widths from
$35425\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ with a median of
$110\text{\,}\mathrm{\SIUnitSymbolAngstrom}$. In the two-line sample, the
faintest source has $V=24.4$. The corresponding H$\alpha$ net flux is
$9.3\text{\times}{10}^{-17}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$,
which at the distance of M101 corresponds to a SFR of
$2.8\text{\times}{10}^{-6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$.
The two-line sample’s H$\alpha$ equivalent widths range across
$8300\text{\,}\mathrm{\SIUnitSymbolAngstrom}$, with a median of
$60\text{\,}\mathrm{\SIUnitSymbolAngstrom}$.

Figures 7-9 show images of the detected sources and their surrounding
environments. Source identifiers for each object were assigned based on their
proximity to known objects in the survey area and their sample (see Figures
7-9). For instance, sources labeled “M101-3-#” are sources near M101 in the
three-line sample. Additional sources were identified as belonging to NGC 5474
and NGC 5486. While the object associated with NGC 5486 lies just inside the
galaxy’s $2R_{25}$ isophotal radius, visual inspection shows the object to be
a distinct source located well outside the galaxy’s star-forming disk, so we
have kept it in our sample.

In the three-line sample, 18 ($95\text{\,}\mathrm{\char 37\relax}$) sources
appear to be associated with the outer disk of M101 and 1
($5\text{\,}\mathrm{\char 37\relax}$) appears to be associated with NGC 5474.
In the two-line sample, 7 ($88\text{\,}\mathrm{\char 37\relax}$) sources
appear to be associated with the outer disk of M101 and 1
($12\text{\,}\mathrm{\char 37\relax}$) are associated with NGC 5486.

###  4.1 Structural Analysis

Since we are looking for faint star-forming objects in the M101 Group, we
begin by comparing the observed properties of our samples to the observed
properties of known satellite galaxies in the Local Universe shifted to the
M101 distance. We focus first on observed size, surface brightness, and
apparent magnitude of our objects in the three- and two-line samples.

Given the better resolution of SDSS, we assigned the sizes of our objects by
utilizing the SDSS structural properties where able. For sources unresolved in
SDSS imaging, we assigned a size equivalent to the FWHM of the SDSS $g$-band
PSF ($None$; Fukugita et al. 1996). For resolved sources, we use the effective
radius of the best-fit de Vaucouleurs or exponential profile as reported in
the SDSS catalog. For objects not in SDSS, we measure the half-light radius
directly from aperture photometry on the $V$-band Schmidt imaging. In this
case, unresolved objects are again given an effective radius equivalent to the
FWHM of the Schmidt imaging ($None$) which corresponds to
$75\text{\,}\mathrm{pc}$ at M101.

We plot these quantities in the structural plots shown in Figure 10. The
unresolved objects that have sizes equivalent to the SDSS or Schmidt PSFs are
marked with arrows; these sizes only serve as upper limits to their true
sizes. For context, the Local Group dwarf galaxies from McConnachie (2012) are
also plotted. Effective radii for the LMC and SMC were taken from Gallart et
al. (2004) and Massana et al. (2020), respectively, while apparent $V$-band
magnitudes were taken from de Vaucouleurs et al. (1991).

Figure 10: Structural properties of our detected sources, compared to Local
Group dwarfs shifted to the M101 distance. Top: Effective radius versus
apparent magnitude. The dotted line indicates $None$ ($33\text{\,}\mathrm{pc}$
at M101’s distance). Bottom: Surface brightness versus apparent magnitude. The
dotted line represents the limit on surface brightness for unresolved sources.
Both the three- and two-line samples are plotted (orange and green markers,
respectively). Pentagons denote objects whose properties were derived using
SDSS data, while squares are sources with properties derived from our Burrell
Schmidt imaging. Data for Local Group dwarfs are taken from McConnachie (2012)
and are shown as gray circles, while known M101 satellites are shown as gray
squares (de Vaucouleurs et al., 1991; Taylor et al., 2005; Merritt et al.,
2014; Bennet et al., 2019; Bellazzini et al., 2020). The limits for H II
galaxies are shown as dashed gray lines (Thuan & Martin, 1981).

We also plotted several of the confirmed and candidate satellite galaxies of
M101 according to Carlsten et al. (2019): NGC 5474, NGC 5477, Holmberg IV,
DF1, DF2, DF3, DwA, and Dw9. There is no single paper that catalogs all of the
photometric properties of the M101 satellite system; quantities for each
galaxy were taken from a variety of sources (de Vaucouleurs et al., 1991;
Taylor et al., 2005; Merritt et al., 2014; Bennet et al., 2019; Bellazzini et
al., 2020). Finally, we also show the region where H II galaxies would lie,
using the definition of Thuan & Martin (1981): $-18.0<M_{V}<-13.5$ and sizes
less than $1\text{\,}\mathrm{kpc}$, which at M101 corresponds to an angular
size of $None$.

In general, the three- and two-line samples are separated into two groups:
there is a group of larger objects that roughly follow the trends of known
dwarf galaxies, and another set of smaller objects that do not. These latter
objects are largely unresolved point sources, giving rise to an upper limit on
size and lower limit on surface brightness as shown in Figure 10. Only one
object lies near the region defined by H II galaxies. We explore what this
source might be by marking each object by associated galaxy.

Figure 11: Effective radius versus apparent magnitude for our detected sources
and comparison samples shifted to the M101 distance. Symbols are the same as
in Figure 10, but color-coded by associated object: M101 (blue), NGC 5474
(orange), or NGC 5486 (green). Top panel shows the three-line sample; bottom
panel shows the two-line sample.

Figure 11 shows the three-line and two-line samples’ structural information
individually, color-coded by associated galaxy. In the three-line sample, most
of the sources that appear to be associated with M101 fall along the smaller
edge of the trend defined by dwarf galaxies. However, these do not appear to
be individual galaxies, but rather H II regions on the outskirts of the spiral
arms (see Figure 8). This includes the two bright sources at $V=15.8$ and
$R_{e}=$$$ and $V=17$ and $R_{e}=$$$; they belong to neighboring large H II
complexes directly north of the center of M101 and both are at the end of the
large distorted spiral arm. These properties make these H II regions very
similar to the known giant extragalactic H II regions of M101, such as NGC
5471 (García-Benito et al., 2011).

The source associated with NGC 5474 is the faintest $V$-band source in the
three-line sample. It lies just outside $2R_{25}$ and does not appear to be an
extension of NGC 5474’s spiral arms like the M101-associated sources (see
Figure 9). Its apparent size is very similar to the Local Group dwarf Segue
II, but given that our source is brighter than Segue II by a factor of 30, it
is unlikely that our source is an undiscovered, very small dwarf galaxy.
Rather, it is likely an H II region; its physical size is
$30\text{\,}\mathrm{pc}$, comparable to the size of a Strömgren sphere powered
by an O9 star (Strömgren, 1939; Osterbrock, 1989).

The two-line sample in Figure 11 has a wider variety of sources than the
three-line sample. As before, all of the sources associated with M101 appear
not as dwarf galaxies, but are interspersed throughout M101’s extended spiral
arms and are likely just H II regions in the outer disk. These regions have a
larger spread in luminosity than the three-line sources, but are all only a
few arcseconds in size. Investigating the line ratios and EWs of the two-line
sources explains why these sources are bright enough to be detected in
H$\alpha$, but not in H$\beta$: half of the sources have high Balmer
decrements, indicating they are heavily extincted, while the other half have
such large uncertainties on their H$\beta$ EW that they could reasonably be
absorption signatures caused by a strong stellar continuum.

In the case of the source near NGC 5486, it is unlikely to be a part of the
galaxy’s outer spiral arms. The galaxy itself has relatively regular outer
isophotes in our deep $B$\- and $V$-band images (see Figure 9). Given that we
can resolve the spatially extended nature of the M101 candidates in our
images, but cannot for this source, it is unlikely that this is a star-forming
object in the M101 Group, but likely more distant. NGC 5486 has a redshift of
$z=0.004$ (de Vaucouleurs et al., 1991); coupled with a Virgocentric flow
model of Mould et al. (2000), this gives a distance to the galaxy of
$28.2\text{\,}\mathrm{Mpc}$. If the detected object is at the same distance as
NGC 5486, it would have a size of $\sim$$1.8\text{\,}\mathrm{kpc}$ and
$M_{V}=-13.5$, similar in physical size to the dwarf irregular galaxy WLM and
similar in luminosity to the Sagittarius Dwarf Spheroidal Galaxy (McConnachie,
2012).

###  4.2 Photometric Analysis

Another way of understanding these objects is to compare their star-forming
properties to those of known galaxies. We do this by again comparing the
observed properties of our samples to the observed properties of known
galaxies shifted to the M101 distance. Figure 12 shows the distribution of
H$\alpha$ flux (left axis) and distance-independent equivalent width for the
three- and two-line samples. As expected, the three-line sample as a whole has
higher EWs than the two-line sample, although both span similar ranges in
emission line flux. Other studies have also used EWs to search for
extragalactic H II regions. Werk et al. (2010) reported their emission line
sources, which they called ELdots, spanned a range of EWs of approximately
$20900\text{\,}\mathrm{\SIUnitSymbolAngstrom}$. Greater than
$150\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ the majority of their ELdots were
H$\alpha$-emitting outlying H II regions; 7 ($27\text{\,}\mathrm{\char
37\relax}$) of our sources have H$\alpha$ EWs greater than
$150\text{\,}\mathrm{\SIUnitSymbolAngstrom}$.

Figure 12: The main panel shows H$\alpha$ flux (for our objects) or
H$\alpha$-inferred SFR (for galaxies in the 11HUGS and SINGG samples) versus
H$\alpha$ equivalent width. Our objects are color-coded by sample and use the
left-hand $y$-axis for flux; smaller marker sizes are used for unresolved
objects. The 11HUGS (black points) and SINGG (blue points) samples are plotted
using the right-hand $y$-axis for absolute SFR. The two axes are equivalent
for objects at the $6.9\text{\,}\mathrm{Mpc}$ distance of M101. The upper
panel shows the normalized distribution of H$\alpha$ equivalent widths in the
various samples, the 11HUGS/SINGG samples together (black outlined bars) and
the three- and two-line samples together (light blue bars).

To place our sources in the context of other star-forming galaxies, we also
plot in Figure 12 the H$\alpha$-inferred SFRs (right axis) and EWs for the
$11\text{\,}\mathrm{Mpc}$ H$\alpha$ and Ultraviolet Galaxy Survey (11HUGS, Lee
et al. 2004; Kennicutt et al. 2008) and the Survey for Ionization in Neutral
Gas Galaxies (SINGG, Meurer et al. 2006). The 11HUGS sample is a virtually
complete sample of Local Volume galaxies with SFRs in the range of
$\sim$${10}^{-5}10\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$
and H$\alpha$ EWs spanning $0545\text{\,}\mathrm{\SIUnitSymbolAngstrom}$. The
SINGG sample utilized those H I-selected galaxies in the H I Parkes All Sky
Survey (HIPASS, Meyer et al. 2004); the SINGG sample has SFRs spanning
$0.001214\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$ and
H$\alpha$ EWs from $2.8451\text{\,}\mathrm{\SIUnitSymbolAngstrom}$.

Comparing our sample to the 11HUGS and SINGG samples, we note several
features. On average our sample, if at the M101 distance, would be objects
that have SFRs equivalent to the faintest galaxies in the 11HUGS and SINGG
samples, i.e. less than
${10}^{-3}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$.
Our sample also consists of objects that have higher H$\alpha$ EWs, on
average, than the faint star-forming galaxies in the two galactic samples.
This is not surprising given that our survey is an emission line survey; we
will find strong emitters with high EWs. If any of our sources are bona-fide
dwarfs, they would represent small objects of low star formation rate but high
equivalent width, i.e. weak starbursting objects.

There are a few galaxies in the 11HUGS/SINGG samples that have such
properties. One galaxy, UGCA 92, a dwarf companion to NGC 1569, has a high EW
($96\text{\,}\mathrm{\SIUnitSymbolAngstrom}$; Kennicutt et al. 2008)
indicating recent star formation. The high EW in UGCA 92 may be due to an
interaction with NGC 1569 (Makarova et al., 2012), making it an interesting
comparison to objects in the M101 Group which may have interacted with M101.

An alternative explanation might be that some of these sources are background
objects. In Figure 12, there is a population of 11HUGS/SINGG galaxies with
high EWs and $-2\leq\log(\text{SFR})\leq-1$, three orders of magnitude
brighter than our sources would be at the M101 distance. Moving our sources
out by a factor of 30 in distance would make them commensurate with those
objects in the 11HUGS/SINGG samples. However, at a distance of
$\sim$$200\text{\,}\mathrm{Mpc}$, the emission lines would have been
redshifted out of our filters and they would not show up in our detections.
Thus it is unlikely that these sources are background objects, and given the
mismatch to known star-forming dwarf galaxies as well as the close proximity
of many of our sources to M101, the bulk of these sources are likely to be
extreme outer-disk H II regions.

Figure 13: The main panel shows H$\alpha$ flux versus EW for points in the
three-line sample (squares) and two-line sample (circles), color-coded by
associated galaxy. The left $y$-axis shows measured flux; the right $y$-axis
shows equivalent star formation if the objects were at the M101 distance.
Smaller marker sizes are used for unresolved objects. The upper panel shows
the normalized distribution of H$\alpha$ equivalent widths of the sources.

Figure 13 is similar to Figure 12, but shows our sources color-coded by their
associated object, with the 11HUGS and SINGG samples removed. We notice that
the objects with high EWs appear to be associated with M101. This is
consistent with the H$\alpha$ EWs of H II regions in the arms of M101 being
high; Cedrés & Cepa (2002) found them to range from
$10e4.5\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ with a median of
$1660\text{\,}\mathrm{\SIUnitSymbolAngstrom}$. This supports our assumption
that these sources are indeed associated with M101 as part of its outer spiral
structure rather than being outlying H II regions.

There are three outliers to the cluster of sources belonging to M101. One
source has a very high emission line flux,
$f_{\text{H}\alpha}=$1.4\text{\times}{10}^{-13}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$$,
and moderately high H$\alpha$ EW,
$112\text{\,}\mathrm{\SIUnitSymbolAngstrom}$. The other has still high but
lower values,
$f_{\text{H}\alpha}=$1.8\text{\times}{10}^{-14}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$$
and $\text{EW(H}\alpha)=$42\text{\,}\mathrm{\SIUnitSymbolAngstrom}$$. These
are the same outliers in the top panel of Figure 11 at $(V,R_{e})=(15.8,$$)$
and $(17,$$)$, respectively. The high emission line fluxes supports the
hypothesis that these sources are similar to the other giant H II regions in
M101’s disk. For instance, the H II region NGC 5471 has an H$\alpha$ flux of
$(3.65\pm
0.17)\text{\times}{10}^{-12}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$
(García-Benito et al., 2011). These sources also have equivalent widths
broadly consistent with H II regions in the inner spiral arms, albeit on the
lower end (Cedrés & Cepa, 2002).

Conversely, the other outlier falls within the range of emission line fluxes
defined by the other M101-associated sources ($(1.8\pm
0.1)\text{\times}{10}^{-15}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$),
but at a much lower H$\alpha$ EW ($15.4\pm
3.1\text{\,}\mathrm{\SIUnitSymbolAngstrom}$). This is the same bright source
in the bottom panel of Figure 11 at $V=18.4$ and $R_{e}=$$$. This is located
in the same group of H II complexes as the source described above, but is
found in an area of H$\beta$ absorption, indicating a stronger continuum.
Given the low EW and H$\beta$ absorption, it is likely that this is a somewhat
older H II region than others that lie in our sample.

The source with the lowest H$\alpha$ EW is the source associated with NGC
5486. If at the distance of NGC 5486, it has an H$\alpha$ luminosity of
$8.7\text{\times}{10}^{37}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}$
giving it an SFR of
$4.7\text{\times}{10}^{-4}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$
(Kennicutt & Evans, 2012). It is interesting to note that its H$\alpha$
luminosity and thus SFR are similar to that of NGC 4163, a nearby dwarf
irregular galaxy (Kennicutt et al., 2008). Their H$\alpha$ EWs are similar as
well; compare our source’s $8\pm 2\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ to
NGC 4163’s $8\pm 2\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ EW (Kennicutt et
al., 2008). It is quite likely that this source is a star-forming satellite of
NGC 5486 similar in structure and luminosity to NGC 4163.

Finally, we turn to the source associated with NGC 5474 with the highest EW in
the entire sample. This source has an H$\alpha$ EW of $422\pm
43\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ putting it firmly in the realm of
H$\alpha$-emitting outlying H II regions as defined by Werk et al. (2010). It
also has moderately strong H$\beta$ and [O III] EWs, $71\pm
22\text{\,}\mathrm{\SIUnitSymbolAngstrom}$ and $282\pm
37\text{\,}\mathrm{\SIUnitSymbolAngstrom}$, respectively. It has a projected
separation from NGC 5474 of $\sim$$None$ ($8.8\text{\,}\mathrm{kpc}$). Given
its structural and photometric properties, this source is likely an outlying H
II region associated with NGC 5474, making it the best (and perhaps only) such
candidate discovered in our survey. Followup spectroscopy would be useful to
secure its status as a bona-fide object in the M101 Group.

##  5 Discussion

Overall, across the $6\text{\,}{\mathrm{deg}}^{2}$ field of our survey, we
detect a total of 19 objects in our three-line sample and 8 objects in our
two-line sample. Of these objects, nearly all are found close to M101, aside
from one object associated with NGC 5474 and another source associated with
NGC 5486 (and thus likely in a background object). None of the emission line
sources detected were located far from bright galaxies, arguing against any
significant, ongoing intragroup star-forming objects in the M101 Group, down
to a limiting star formation rate of
$1.7\text{\times}{10}^{-6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$.

In the discussion that follows, we will investigate the consequences of this
rarity of star-forming objects in the context of intragroup H I clouds and
newly discovered ultradiffuse galaxies in the M101 Group, as well as the faint
end of the star-forming luminosity function. We will also give a more broad
discussion of the merger history of M101 and group environments in general.

Our survey area covers the entire extent of the deepest portion of the M101 H
I imaging survey by Mihos et al. (2012). That survey had a limiting H I column
density of $\log(N)=16.8$ and H I mass detection limit of
$2\text{\times}{10}^{6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}$; for comparison,
that survey would have detected even the lowest H I mass objects in the SINGG
survey if they were in the M101 Group. The Mihos et al. (2012) survey did
detect a number of discrete H I clouds in the M101 Group, along with a diffuse
loop of H I extending $85\text{\,}\mathrm{kpc}$ to the southwest of M101. This
loop and the associated H I clouds likely arise from tidal interactions
between M101 and its companions, yet our deep narrowband imaging presented
here shows no evidence of ongoing star formation in this gas, either in
discrete sources or diffuse emission. Nor did the deep broadband imaging of
the M101 system by Mihos et al. (2013) show evidence for diffuse light in this
gas. If extended star formation was triggered in this gas by the past
interactions in the M101 Group, that star formation must have been very weak
and died out quickly.

We have also searched for faint H$\alpha$ emission in several of the recently
discovered ultradiffuse galaxies in the M101 Group. Five of these dwarfs are
seen in our broadband images: DF1, DF2, DF3, DwA, and Dw9 (Merritt et al.,
2014, 2016; Karachentsev et al., 2015; Danieli et al., 2017; Carlsten et al.,
2019). Bennet et al. (2017) used a large NUV footprint and determined that all
five of the dwarfs in our images lack a NUV excess indicating an upper limit
SFR of $(1.7\pm
0.5)\text{\times}{10}^{-3}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$.
We detect no compact emission line sources associated with these objects, and
aperature photometry over the $None$ sizes (typical half-light radii as
reported by Bennet et al. 2017) reveals no diffuse emission down to a level of
$4.5\text{\times}{10}^{35}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}$.
This places a strong upper limit on any star formation in these objects of
$\text{SFR}<$2.5\text{\times}{10}^{-6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$$.

This lack of detected H$\alpha$ emission is consistent with the red optical
colors of these dwarfs (Merritt et al., 2014; Bennet et al., 2017), arguing
that these are older objects, rather than systems formed during recent tidal
interactions in the M101 Group. UDGs are observed in both the field and the
group environments, so an intuitive formation scenario for those in group
environments would be that they were already “puffed up” in the field and
quenched after falling into the group (Román & Trujillo, 2017; Alabi et al.,
2018; Chan et al., 2018; Ferré-Mateu et al., 2018). Recent observations seem
to support this scenario – Román & Trujillo (2017) and Alabi et al. (2018)
found UDGs to have redder colors at smaller cluster-centric distances.
Simulations have predicted that of all satellites to have ever existed in a
group environment, half originated in the field as UDGs (Jiang et al., 2019),
themselves formed from supernovae feedback (Di Cintio et al., 2017), and half
were normal galaxies puffed up by tides as satellites (Jiang et al., 2019).
Perhaps the lopsidedness of M101’s satellites might indicate that they are
part of a infalling low-mass group (Merritt et al., 2014); indeed most
isolated blue galaxies, like M101, have lopsided satellite distributions
(Brainerd & Samuels, 2020). More detailed kinematic investigation is needed of
the M101 Group to determine the origin of these UDGs.

Figure 14: The ratio of the expected number of faint-to-bright objects in the
M101 Group as a function of the luminosity faint end slope, $\alpha$. Bright
objects are defined as having SFRs of
$$0.01$<$\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$<$50$$, while
fainter ones have SFRs of
$${10}^{-5}$<$\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$<$0.01$$.
The gray shaded area shows where the faint sources are rare, as indicated by
our data, and favors slopes flatter than $\alpha\sim-1.0$.

Given the lack of a significant population of low luminosity star-forming
dwarf galaxies in the M101 Group, we can place rough limits on the slope of
the faint end of the star-forming luminosity function of galaxies in the
group. A steep faint end slope would predict a high dwarf-to-giant ratio; the
lack of these dwarfs in the M101 Group argues instead for a relatively shallow
slope. By adopting a Schechter (1976) function with
$\text{SFR}_{\ast}=$9\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$$
(Bothwell et al., 2011) and varying the faint end slope, we can calculate how
many faint star-forming objects we should detect within the group.

In the M101 Group, there are four galaxies with SFRs greater than
$0.01\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$: M101,
NGC 5474, NGC 5477, and Holmberg IV (Kennicutt et al., 2008; Kennicutt &
Evans, 2012). However, our deep, wide-field imaging detects no intragroup
objects with lower star formation rates. While it is possible that some
objects may have been missed due to, for example, being projected directly in
front or behind M101’s disk, our results argue that the number of faint star
forming galaxies (down to
$\text{SFRs}\sim${10}^{-5}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$$)
must be very low, and certainly do not outnumber the brighter galaxies. This
is contrary to expectation if the faint end of the star-forming luminosity
function is even moderately steep. For example, as illustrated in Figure 14, a
flat luminosity function (with slope $\alpha=-1$) would predict equal numbers
of objects in the low star formation range of
$${10}^{-5}$<\text{SFR}<${10}^{-2}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$$,
while a steeper slope with $\alpha=-1.5$ predicts on the order of 120 objects
in the same SFR range. Since we found no evidence of such a population, either
flatter values for the faint end slope ($\alpha>-1.0$), or a luminosity
function that is sharply truncated below
$0.01\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$, are
required.

Most of the sources detected ($92\text{\,}\mathrm{\char 37\relax}$) were
associated with M101 and are not true outlying H II regions as they were
extensions of the spiral arms. However, one source located near NGC 5474 had
both photometric and structural properties consistent with previous searches
for isolated H II regions. We can quantify this source by calculating the
total number of ionizing photons, $Q_{0}$, from Osterbrock (1989):

$Q_{0}=2.2\frac{L_{\text{H}\alpha}}{h\nu_{\text{H}\alpha}}\sim$7.2\text{\times}{10}^{11}$L_{\text{H}\alpha}\,${\mathrm{s}}^{-1}$,$
(1)

where we assume case B recombination (Osterbrock, 1989). Assuming this source
to be at the M101 distance, it could be powered by only 4 O9V stars (Martins
et al., 2005), similar to the faintest isolated H II regions detected in the
studies of other galaxies by Ryan-Weber et al. (2004) and Werk et al. (2010).
Such a low level of star formation pushes the ill-defined boundary
differentiating outlying H II regions from star-forming dwarf galaxies;
follow-up spectroscopy would be of interest in studying this source in more
detail.

Another source that deserves spectroscopic follow-up is N5486-2-1. Although
this was located within $2R_{25}$ of NGC 5486 and does not satisfy the
classical definition of an outlying H II region, assuming it is physically
near NGC 5486, it displays many properties similar to local dwarf galaxies.
The knot of star formation our survey targeted could be powered by $\sim$80
O9V stars (Martins et al., 2005), consistent with dwarf galaxies (Werk et al.,
2010). Additionally, this source has a similar $V$-band and H$\alpha$
luminosity to the dwarf irregular galaxy NGC 4163 as mentioned above. Another
striking resemblance is the asymmetry in the broadband images. The source
appears to consist of a bright compact region with a more diffuse extension to
the west, structure which is also evident in the shallower $gri$ composite
from SDSS. Similarly NGC 4163 is itself asymmetric in nature, with a burst of
star formation occurring in the central portion of the galaxy, but not the
outer regions (McQuinn et al., 2012). It also has a peculiar H I distribution,
with an H I tail to the west and possibly the south (Hunter et al., 2011;
Lelli et al., 2014). Although it is not clear what could be causing the
behavior in NGC 4163, perhaps some large-scale interaction with NGC 5486 has
caused this asymmetry in our source.

Equally uncertain is the interaction history of the M101 Group as a whole. The
asymmetric disk of M101 has long been believed to arise from an interaction
(Beale & Davies, 1969; Rownd et al., 1994; Waller et al., 1997). Low surface
brightness optical light has been detected in the outskirts of M101 (Mihos et
al., 2013) with colors and stellar populations consistent with a burst of star
formation $\sim$$300400\text{\,}\mathrm{Myr}$ ago (Mihos et al., 2018). High-
velocity H I gas has also been observed in the same location as the optical
light (van der Hulst & Sancisi, 1988; Mihos et al., 2012), while intermediate-
velocity H I gas has been detected between M101 and NGC 5474 (Mihos et al.,
2012). NGC 5474’s offset bulge is usually added as further evidence of an
interaction, although recent work has called that into question (Bellazzini et
al., 2020; Pascale et al., 2021).

Galaxy-galaxy interactions frequently give rise to star formation in tidal
debris (Schombert et al., 1990; Gerhard et al., 2002; Sakai et al., 2002;
Cortese et al., 2004; Ryan-Weber et al., 2004; Boquien et al., 2007, 2009;
Werk et al., 2010). Intragroup H II regions have also been detected between
pairs of massive galaxies in compact groups (Sakai et al., 2002; Mendes de
Oliveira et al., 2004), illustrating that the small group environment easily
drives interactions leading to star-forming regions beyond the galactic disks.

Given our lack of detection of any isolated, intragroup H II regions, what
does this mean for the interaction scenario? All of the examples of star-
forming objects in interacting systems involve strong interactions or major-
merger events. The lack of star-forming H II regions between NGC 5474 and M101
might indicate that the two galaxies have undergone only a weak interaction,
or that any tidally-triggered intragroup star formation was very short-lived.
If the pair’s luminosity ratio of $17:1$ in the $V$-band is indicative of the
masses of the galaxies, then this would classify it as a low-mass encounter.

Furthermore, the M101 Group is not a compact galaxy group. Compact galaxy
groups are very dense environments, containing only a few galaxies separated
by distances comparable to their sizes. This makes them strongly interacting
environments, leading to the formation of tidal tails, intragroup star
formation, and tidal dwarf galaxies (e.g. de Mello et al., 2008; Torres-Flores
et al., 2009). A well-known example of this is Stephan’s Quintet (Arp 319)
where the star formation is heavily influenced by the ongoing interactions
between its member galaxies. Intragroup star formation can be found at the tip
of a shock front (Xu et al., 1999, 2003) and in a tidal tail (Arp, 1973;
Sulentic et al., 2001; Xu et al., 2005). Numerous tidal dwarf galaxy
candidates have been found in tidal tails as well (Hunsberger et al., 1996).

In contrast with Stephan’s Quintet, the M101 Group is dominated by M101, with
relatively few low-mass companions, making it possibly the poorest group in
the Local Volume (Bremnes et al., 1999). The lack of an abundance of
intragroup star formation in the M101 Group might be because the M101 Group
involves only weak interactions with low mass companions. Given M101’s
“anemic” stellar halo (van Dokkum et al., 2014; Jang et al., 2020), it is
unlikely that M101 has gone through any major mergers, and very few, if any,
minor mergers. With no comparable companions nearby, it is likely that no
outlying H II regions will form until one of the low mass companions falls
into and merges with M101. Clearly, detailed computer modeling is needed to
unravel all of the features of the M101 Group.

##  6 Summary

We have conducted a narrowband emission line survey of the M101 Group to
search for faint star-forming dwarf galaxies and outlying H II regions. Using
narrowband filters that target H$\alpha$, H$\beta$, and [O III], our survey
detects sources as faint as
$f_{\text{H}\alpha}=$5.7\text{\times}{10}^{-17}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$$,
or an equivalent star formation rate of
$1.7\text{\times}{10}^{-6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$
for sources in the M101 Group. Imaging in multiple lines significantly reduces
contamination from single-line detections of redshifted sources in the
background, but there is significant contamination from Milky Way stars,
particularly M stars whose complicated continuum structure mimics emission
line flux in our narrowband filters. To eliminate this contamination, we use
stellar spectral templates to place equivalent width cuts on the sample,
requiring
$\text{EW}(\text{H}\alpha)>$8\text{\,}\mathrm{\SIUnitSymbolAngstrom}$$,
$\text{EW}(\text{H}\beta)>$2\text{\,}\mathrm{\SIUnitSymbolAngstrom}$$, and
$\text{EW}(\text{[{O~{}III}]})>$5\text{\,}\mathrm{\SIUnitSymbolAngstrom}$$. We
further cross-match the sample with the _Gaia_ EDR3 catalog to eliminate
objects with detected parallax and proper motion. Finally, we cross-match the
remaining sources with the SDSS catalog to remove known background sources
and, in cases where H$\beta$ is detected, we also reject sources with
anomalous (unphysical) Balmer decrements. Our final sample consists of 17
sources detected in all three of the H$\alpha$, H$\beta$, and [O III] filters
(the three-line sample), and another 8 sources detected only in H$\alpha$ and
[O III] (the two-line sample).

  1. 1.

Of the 27 total sources, 25 ($93\text{\,}\mathrm{\char 37\relax}$) were
located in M101’s extreme outer disk, while two were associated with other
bright galaxies in the survey area. We found no evidence for any truly
outlying emission line regions in the M101 intragroup environment.

  2. 2.

Assuming that our sources are at the M101 distance, a comparison between their
physical properties and the properties of known Local Group dwarfs and star-
forming galaxies in the 11HUGS and SINGG surveys argues that most of our
sources are not bona-fide star-forming dwarfs. Instead, their small sizes, low
H$\alpha$ luminosities, and high H$\alpha$ EWs argue that they are more likely
M101 outer disk H II regions.

  3. 3.

Two detected sources were unassociated with M101 itself: an unresolved but
high EW source near the M101 satellite NGC 5474 (likely an outlying H II
region), and a spatially extended, low EW satellite of the background galaxy
NGC 5486.

  4. 4.

We searched for ionized gas emission in five previously discovered
ultradiffuse dwarf galaxy candidates in the M101 Group (DF1, DF2, DF3, DwA,
Dw9) and found none to a limit of
$2.5\text{\times}{10}^{-6}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$.
This makes it unlikely that any of these ultradiffuse candidates are
undergoing star formation at present times.

  5. 5.

The lack of any significant population of star-forming objects in the M101
Group down to a limiting equivalent star formation rate of
${10}^{-5}\text{\,}\mathrm{\mbox{$M_{\odot}$}}\text{\,}{\mathrm{yr}}^{-1}$
argues for a shallow faint end slope of the star-forming luminosity function
($\alpha\sim 1$).

  6. 6.

Given that tidally induced star formation is a common outcome of close galaxy
encounters, the lack of outlying H II regions also suggests either that
interactions within the M101 Group have been relatively weak, or that any
extended star formation triggered by the encounters has quickly died out. Both
scenarios are consistent with recent studies attributing M101’s distorted
morphology to a weak, retrograde encounter with NGC 5474 some
$300400\text{\,}\mathrm{Myr}$ ago.

This work was supported in part by a Towson Memorial Scholarship to RG, and by
grants to JCM from the National Science Foundation (award 1108964) and the Mt
Cuba Astronomical Foundation. We also thank Stacy McGaugh for helpful
discussions. =0.7in

Table 3: Narrowband Properties

ID | RA | Dec | E/U |  $\Delta f_{\text{H}\alpha}$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000). | $\Delta f_{\text{H}\beta}$ | $\Delta f_{\text{[{O~{}III}]}}$ | EW${}_{\text{H}\alpha}$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000). | EW${}_{\text{H}\beta}$ | EW${}_{\text{[{O~{}III}]}}$ |  $\text{H}\alpha/\text{H}\beta$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000). |  $\text{[{O~{}III}]}/\text{H}\alpha$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000).
---|---|---|---|---|---|---|---|---|---|---|---
| [Deg] | [Deg] |  | [${10}^{-16}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$] | [${10}^{-16}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$] | [${10}^{-16}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$] | [$\mathrm{\SIUnitSymbolAngstrom}$] | [$\mathrm{\SIUnitSymbolAngstrom}$] | [$\mathrm{\SIUnitSymbolAngstrom}$] |  |
M101-3-1 | 210.7923 | 54.5976 | E | 51.167 (1.637) | 19.265 (2.275) | 20.744 (2.190) | 83.947 (16.911) | 22.162 (6.247) | 26.814 (6.174) | 2.656 | 0.405
M101-3-2 | 210.8047 | 54.5948 | E | 1446.379 (3.644) | 520.204 (4.936) | 1827.889 (4.853) | 111.643 (48.094) | 25.178 (13.898) | 102.963 (45.658) | 2.780 | 1.264
M101-3-3 | 210.8159 | 54.5954 | E | 42.782 (1.494) | 8.897 (2.075) | 19.030 (2.001) | 35.695 (6.600) | 4.523 (1.493) | 11.347 (2.435) | 4.809 | 0.445
M101-3-4 | 210.8268 | 54.5884 | E | 179.646 (2.756) | 28.383 (3.829) | 172.745 (3.704) | 42.363 (14.186) | 3.854 (1.740) | 26.896 (9.307) | 6.329 | 0.962
M101-3-5 | 210.8277 | 54.5944 | E | 48.542 (1.654) | 13.523 (2.300) | 40.987 (2.218) | 49.854 (10.167) | 8.679 (2.687) | 30.387 (6.503) | 3.590 | 0.844
M101-3-6 | 210.9173 | 54.6299 | U | 9.192 (0.629) | 4.339 (0.874) | 8.228 (0.842) | 287.044 (29.440) | 254.578 (57.042) | 511.041 (65.872) | 2.118 | 0.895
M101-3-7 | 211.0488 | 54.5866 | U | 14.315 (0.858) | 6.166 (1.192) | 9.960 (1.148) | 106.914 (12.851) | 37.397 (8.795) | 63.710 (10.026) | 2.322 | 0.696
M101-3-8 | 211.0629 | 54.5992 | E | 5.951 (0.558) | 3.063 (0.774) | 2.935 (0.745) | 100.021 (11.582) | 87.032 (23.253) | 61.645 (16.232) | 1.943 | 0.493
M101-3-9 | 211.1460 | 54.4876 | U | 11.338 (0.848) | 6.481 (1.178) | 6.802 (1.133) | 108.648 (13.840) | 73.776 (16.600) | 87.410 (17.255) | 1.749 | 0.600
M101-3-10 | 211.1545 | 54.4514 | U | 18.797 (1.105) | 5.446 (1.537) | 12.977 (1.480) | 51.262 (7.518) | 9.571 (3.168) | 28.209 (5.053) | 3.452 | 0.690
M101-3-11 | 211.1651 | 54.4407 | U | 26.721 (0.961) | 10.942 (1.331) | 31.493 (1.285) | 104.887 (12.783) | 32.559 (6.279) | 113.648 (14.358) | 2.442 | 1.179
M101-3-12 | 211.1661 | 54.4580 | U | 54.601 (1.483) | 22.772 (2.056) | 107.590 (1.990) | 144.730 (26.330) | 58.293 (14.470) | 360.163 (66.812) | 2.398 | 1.970
M101-3-13 | 211.1676 | 54.2532 | E | 49.412 (1.249) | 22.497 (1.728) | 112.280 (1.678) | 159.882 (24.508) | 76.696 (16.012) | 426.537 (66.469) | 2.196 | 2.272
M101-3-14 | 211.1932 | 54.4407 | E | 10.295 (0.929) | 4.341 (1.290) | 7.437 (1.243) | 44.510 (6.442) | 17.032 (5.635) | 31.805 (6.473) | 2.372 | 0.722
M101-3-15 | 211.1955 | 54.3812 | E | 9.888 (0.917) | 4.552 (1.278) | 4.716 (1.228) | 77.771 (11.291) | 28.326 (8.932) | 36.478 (10.382) | 2.172 | 0.477
M101-3-16 | 211.1994 | 54.4902 | U | 11.352 (0.755) | 5.820 (1.050) | 3.253 (1.009) | 156.237 (17.710) | 79.776 (17.197) | 49.120 (15.925) | 1.950 | 0.287
M101-3-17 | 211.2051 | 54.3800 | E | 30.155 (1.081) | 13.725 (1.498) | 9.856 (1.439) | 111.209 (15.120) | 49.666 (9.963) | 46.681 (9.261) | 2.197 | 0.327
M101-3-18 | 211.2138 | 54.3906 | E | 12.376 (0.755) | 5.801 (1.049) | 11.343 (1.011) | 205.514 (22.631) | 90.213 (19.485) | 204.745 (26.554) | 2.133 | 0.917
N5474-3-1 | 211.2695 | 53.7356 | U | 8.000 (0.578) | 2.646 (0.804) | 7.130 (0.774) | 422.339 (42.488) | 70.703 (22.404) | 281.790 (36.729) | 3.023 | 0.891

Note. — Numbers in parentheses are the associated uncertainties for that
quantity.

=0.7in

Table 4: Narrowband Properties

ID | RA | Dec | E/U |  $\Delta f_{\text{H}\alpha}$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000). | $\Delta f_{\text{H}\beta}$ | $\Delta f_{\text{[{O~{}III}]}}$ | EW${}_{\text{H}\alpha}$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000). | EW${}_{\text{H}\beta}$ | EW${}_{\text{[{O~{}III}]}}$ |  $\text{H}\alpha/\text{H}\beta$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000). |  $\text{[{O~{}III}]}/\text{H}\alpha$aaH$\alpha$ fluxes and EWs have been corrected for [N II] emission by assuming [N II]/H$\alpha=0.33$ (Kennicutt, 1992; Jansen et al., 2000).
---|---|---|---|---|---|---|---|---|---|---|---
| [Deg] | [Deg] |  | [${10}^{-16}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$] | [${10}^{-16}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$] | [${10}^{-16}\text{\,}\mathrm{erg}\text{\,}{\mathrm{s}}^{-1}\text{\,}{\mathrm{cm}}^{-2}$] | [$\mathrm{\SIUnitSymbolAngstrom}$] | [$\mathrm{\SIUnitSymbolAngstrom}$] | [$\mathrm{\SIUnitSymbolAngstrom}$] |  |
M101-2-1 | 210.8156 | 54.5862 | E | 18.049 (1.500) | -9.414 (2.092) | 29.587 (2.025) | 15.397 (3.094) | -4.529 (-1.467) | 15.829 (3.183) | -1.917 | 1.639
M101-2-2 | 210.9127 | 54.5898 | E | 1.596 (0.457) | 0.374 (0.638) | 3.050 (0.614) | 60.755 (17.741) | 11.030 (18.830) | 125.374 (26.252) | 4.267 | 1.911
M101-2-3 | 211.1586 | 54.2476 | E | 5.978 (0.824) | 1.448 (1.150) | 6.867 (1.107) | 59.436 (10.137) | 8.661 (6.968) | 48.588 (9.309) | 4.128 | 1.149
M101-2-4 | 211.1632 | 54.4533 | E | 4.724 (0.684) | 0.524 (0.955) | 5.957 (0.919) | 80.791 (13.505) | 5.584 (10.199) | 79.426 (14.028) | 9.022 | 1.261
M101-2-5 | 211.1730 | 54.4580 | E | 0.929 (0.248) | 0.347 (0.345) | 1.049 (0.333) | 295.520 (79.335) | 117.886 (117.421) | 325.623 (103.728) | 2.679 | 1.129
M101-2-6 | 211.2291 | 54.3049 | U | 1.436 (0.386) | 0.570 (0.537) | 1.782 (0.517) | 49.596 (13.533) | 17.634 (16.640) | 70.720 (20.794) | 2.518 | 1.241
M101-2-7 | 211.2296 | 54.2999 | E | 2.764 (0.369) | 1.219 (0.513) | 3.219 (0.494) | 249.438 (35.095) | 97.326 (41.319) | 265.051 (42.508) | 2.267 | 1.165
N5486-2-1 | 211.8695 | 55.0837 | E | 9.093 (1.306) | 3.388 (1.800) | 8.630 (1.741) | 8.253 (1.771) | 3.254 (1.849) | 7.853 (2.036) | 2.683 | 0.949

Note. — Numbers in parentheses are the associated uncertainties for that
quantity.

Table 5: Three-Line Sample Broadband Properties ID | $V$ | $B-V$ | $u$ | $g$ | $r$ | $i$ | $z$ | $FUV$ | $NUV$
---|---|---|---|---|---|---|---|---|---
M101-3-1 | 19.20 | 0.05 | – | – | – | – | – | – | –
M101-3-2 | 15.80 | 0.10 | 21.50 | 21.60 | 21.39 | 21.10 | 20.55 | – | –
M101-3-3 | 18.45 | 0.05 | – | – | – | – | – | – | –
M101-3-4 | 16.98 | 0.09 | 18.71 | 18.56 | 18.65 | 18.92 | 22.70 | – | –
M101-3-5 | 18.61 | 0.10 | – | – | – | – | – | – | –
M101-3-6 | 22.46 | 0.28 | 23.67 | 22.49 | 23.36 | 23.79 | 23.54 | – | –
M101-3-7 | 20.83 | 0.06 | 22.47 | 22.33 | 22.23 | 24.29 | 23.24 | – | –
M101-3-8 | 21.91 | 0.32 | – | – | – | – | – | – | –
M101-3-9 | 21.31 | -0.01 | 21.33 | 21.50 | 21.64 | 22.19 | 22.61 | 17.55 | 18.87
M101-3-10 | 19.74 | 0.05 | 20.72 | 19.93 | 20.12 | 21.25 | 21.36 | 16.99 | 18.14
M101-3-11 | 20.12 | 0.12 | 21.07 | 21.13 | 21.23 | 22.02 | 22.18 | 17.23 | 18.40
M101-3-12 | 19.66 | 0.27 | 21.57 | 20.73 | 20.66 | 22.35 | 21.47 | 16.28 | 17.47
M101-3-13 | 19.80 | 0.59 | 19.79 | 19.71 | 19.78 | 20.81 | 19.24 | 16.48 | 17.96
M101-3-14 | 20.34 | 0.41 | – | – | – | – | – | 17.05 | 18.34
M101-3-15 | 20.92 | 0.10 | – | – | – | – | – | 16.96 | 18.35
M101-3-16 | 21.55 | 0.15 | 22.46 | 22.14 | 23.03 | 22.84 | 21.40 | 17.49 | 18.73
M101-3-17 | 20.39 | 0.08 | – | – | – | – | – | 16.56 | 17.94
M101-3-18 | 21.49 | 0.26 | – | – | – | – | – | 17.30 | 18.69
N5474-3-1 | 22.89 | -0.16 | 22.83 | 22.46 | 22.77 | 23.74 | 23.61 | – | –
Table 6: Two-Line Sample Broadband Properties ID | $V$ | $B-V$ | $u$ | $g$ | $r$ | $i$ | $z$ | $FUV$ | $NUV$
---|---|---|---|---|---|---|---|---|---
M101-2-1 | 18.37 | 0.08 | – | – | – | – | – | – | –
M101-2-2 | 22.67 | 0.07 | – | – | – | – | – | – | –
M101-2-3 | 21.01 | 0.10 | – | – | – | – | – | 17.39 | 18.86
M101-2-4 | 21.45 | 0.21 | – | – | – | – | – | 17.98 | 19.22
M101-2-5 | 24.37 | 0.25 | – | – | – | – | – | 20.12 | 21.36
M101-2-6 | 22.75 | 0.35 | 23.35 | 22.56 | 22.34 | 22.44 | 22.42 | 18.73 | 20.28
M101-2-7 | 23.24 | 0.31 | – | – | – | – | – | 18.83 | 20.38
N5486-2-1 | 18.79 | 0.55 | 19.51 | 19.17 | 19.24 | 19.14 | 23.58 | – | –

## References

  * Ahumada et al. (2020) Ahumada, R., Allende Prieto, C., Almeida, A., et al. 2020, ApJS, 249, 3
  * Alabi et al. (2018) Alabi, A., Ferré-Mateu, A., Romanowsky, A. J., et al. 2018, MNRAS, 479, 3308
  * Ann et al. (2015) Ann, H. B., Seo, M., & Ha, D. K. 2015, ApJS, 217, 27
  * Arp (1973) Arp, H. 1973, ApJ, 183, 411
  * Astropy Collaboration et al. (2013) Astropy Collaboration, Robitaille, T. P., Tollerud, E. J., et al. 2013, A&A, 558, A33
  * Astropy Collaboration et al. (2018) Astropy Collaboration, Price-Whelan, A. M., Sipőcz, B. M., et al. 2018, AJ, 156, 123
  * Beale & Davies (1969) Beale, J. S., & Davies, R. D. 1969, Natur, 221, 531
  * Bellazzini et al. (2020) Bellazzini, M., Annibali, F., Tosi, M., et al. 2020, A&A, 634, A124
  * Bennet et al. (2019) Bennet, P., Sand, D. J., Crnojević, D., et al. 2019, ApJ, 885, 153
  * Bennet et al. (2017) —. 2017, ApJ, 850, 109
  * Bertin & Arnouts (1996) Bertin, E., & Arnouts, S. 1996, A&AS, 117, 393
  * Boquien et al. (2007) Boquien, M., Duc, P.-A., Braine, J., et al. 2007, A&A, 467, 93
  * Boquien et al. (2009) Boquien, M., Duc, P.-A., Wu, Y., et al. 2009, AJ, 137, 4561
  * Bothwell et al. (2011) Bothwell, M. S., Kennicutt, R. C., Johnson, B. D., et al. 2011, MNRAS, 415, 1815
  * Bradley et al. (2019) Bradley, L., Sipőcz, B., Robitaille, T., et al. 2019, astropy/photutils: v0.7.2
  * Brainerd & Samuels (2020) Brainerd, T. G., & Samuels, A. 2020, ApJ, 898, L15
  * Bremnes et al. (1999) Bremnes, T., Binggeli, B., & Prugniel, P. 1999, A&AS, 137, 337
  * Brown et al. (2014) Brown, T. M., Tumlinson, J., Geha, M., et al. 2014, ApJ, 796, 91
  * Cannon et al. (2011) Cannon, J. M., Giovanelli, R., Haynes, M. P., et al. 2011, ApJ, 739, L22
  * Cannon et al. (2018) Cannon, J. M., Shen, Z., McQuinn, K. B. W., et al. 2018, ApJ, 864, L14
  * Carlsten et al. (2019) Carlsten, S. G., Beaton, R. L., Greco, J. P., & Greene, J. E. 2019, ApJ, 878, L16
  * Cedrés & Cepa (2002) Cedrés, B., & Cepa, J. 2002, A&A, 391, 809
  * Chan et al. (2018) Chan, T. K., Kereš, D., Wetzel, A., et al. 2018, MNRAS, 478, 906
  * Cortese et al. (2004) Cortese, L., Gavazzi, G., Boselli, A., & Iglesias-Paramo, J. 2004, A&A, 416, 119
  * Danieli et al. (2017) Danieli, S., van Dokkum, P., Merritt, A., et al. 2017, ApJ, 837, 136
  * de Lapparent (2003) de Lapparent, V. 2003, A&A, 408, 845
  * de Mello et al. (2008) de Mello, D. F., Torres-Flores, S., & Mendes de Oliveira, C. 2008, AJ, 135, 319
  * de Vaucouleurs et al. (1991) de Vaucouleurs, G., de Vaucouleurs, A., Corwin, H. G., et al. 1991, Third Reference Catalogue of Bright Galaxies (Springer)
  * Di Cintio et al. (2017) Di Cintio, A., Brook, C. B., Dutton, A. A., et al. 2017, MNRAS, 466, L1
  * Domínguez et al. (2013) Domínguez, A., Siana, B., Henry, A. L., et al. 2013, ApJ, 763, 145
  * Ferguson et al. (1998) Ferguson, A. M. N., Wyse, R. F. G., Gallagher, J. S., & Hunter, D. A. 1998, ApJ, 506, L19
  * Ferré-Mateu et al. (2018) Ferré-Mateu, A., Alabi, A., Forbes, D. A., et al. 2018, MNRAS, 479, 4891
  * Fukugita et al. (1996) Fukugita, M., Ichikawa, T., Gunn, J. E., et al. 1996, AJ, 111, 1748
  * Gaia Collaboration et al. (2020) Gaia Collaboration, Brown, A. G. A., Vallenari, A., et al. 2020, arXiv:2012.01533
  * Gallagher & Hunter (1984) Gallagher, J. S., & Hunter, D. A. 1984, ARA&A, 22, 37
  * Gallart et al. (2004) Gallart, C., Stetson, P. B., Hardy, E., Pont, F., & Zinn, R. 2004, ApJ, 614, L109
  * García-Benito et al. (2011) García-Benito, R., Pérez, E., Díaz, A. I., Maíz Apellániz, J., & no, M. C. 2011, AJ, 141, 126
  * Gerhard et al. (2002) Gerhard, O., Arnaboldi, M., Freeman, K. C., & Okamura, S. 2002, ApJ, 580, L121
  * Harris et al. (2020) Harris, C. R., Millman, K. J., van der Walt, S. J., et al. 2020, Natur, 585, 357
  * Hodge (1975) Hodge, P. W. 1975, ApJ, 201, 556
  * Høg et al. (2000) Høg, E., Fabricius, C., Makarov, V. V., et al. 2000, A&A, 355, L27
  * Hunsberger et al. (1996) Hunsberger, S. D., Charlton, J. C., & Zaritsky, D. 1996, ApJ, 462, 50
  * Hunter et al. (1998) Hunter, D. A., Elmegreen, B. G., & Baker, A. L. 1998, ApJ, 493, 595
  * Hunter et al. (2011) Hunter, D. A., Elmegreen, B. G., Oh, S.-H., et al. 2011, AJ, 143, 121
  * Hunter (2007) Hunter, J. D. 2007, CSE, 9, 90
  * Jang et al. (2020) Jang, I. S., de Jong, R. S., Holwerda, B. W., et al. 2020, A&A, 637, A8
  * Jansen et al. (2000) Jansen, R. A., Fabricant, D., Franx, M., & Caldwell, N. 2000, ApJS, 126, 331
  * Jiang et al. (2019) Jiang, F., Dekel, A., Freundlich, J., et al. 2019, MNRAS, 487, 5272
  * Karachentsev et al. (2015) Karachentsev, I. D., Riepe, P., Zilch, T., et al. 2015, AstBu, 70, 379
  * Keel et al. (2012) Keel, W. C., Lintott, C. J., Schawinski, K., et al. 2012, AJ, 144, 66
  * Kellar et al. (2012) Kellar, J. A., Salzer, J. J., Wegner, G., Gronwall, C., & Williams, A. 2012, AJ, 143, 145
  * Kennicutt (1989) Kennicutt, R. C. 1989, ApJ, 344, 685
  * Kennicutt (1992) —. 1992, ApJ, 388, 310
  * Kennicutt (1998) —. 1998, ARA&A, 36, 189
  * Kennicutt & Evans (2012) Kennicutt, R. C., & Evans, N. J. 2012, ARA&A, 50, 531
  * Kennicutt et al. (2008) Kennicutt, R. C., Lee, J. C., Funes, J. G., et al. 2008, ApJS, 178, 247
  * Lee et al. (2007) Lee, J. C., Kennicutt, R. C., Funes, J. G., Sakai, S., & Akiyama, S. 2007, ApJ, 671, L113
  * Lee et al. (2009) —. 2009, ApJ, 692, 1305
  * Lee et al. (2004) Lee, J. C., Kennicutt, R. C., Funes, J. G., et al. 2004, BAAS, 36, 1442
  * Lelièvre & Roy (2000) Lelièvre, M., & Roy, J.-R. 2000, AJ, 120, 1306
  * Lelli et al. (2014) Lelli, F., Verheijen, M., & Fraternali, F. 2014, MNRAS, 445, 1694
  * Lotz et al. (2004) Lotz, J. M., Primack, J., & Madau, P. 2004, AJ, 128, 163
  * Lupton et al. (1999) Lupton, R. H., Gunn, J. E., & Szalay, A. S. 1999, AJ, 118, 1406
  * Makarova et al. (2012) Makarova, L., Makarov, D., & Savchenko, S. 2012, MNRAS, 423, 294
  * Martin & Kennicutt (2001) Martin, C. L., & Kennicutt, R. C. 2001, ApJ, 555, 301
  * Martins et al. (2005) Martins, F., Schaerer, D., & Hillier, D. J. 2005, A&A, 436, 1049
  * Massana et al. (2020) Massana, P., Noël, N. E. D., Nidever, D. L., et al. 2020, MNRAS, 498, 1034
  * Massey et al. (1988) Massey, P., Strobel, K., Barnes, J. V., & Anderson, E. 1988, ApJ, 328, 315
  * Matheson et al. (2012) Matheson, T., Joyce, R. R., Allen, L. E., et al. 2012, ApJ, 754, 19
  * McConnachie (2012) McConnachie, A. W. 2012, AJ, 144, 4
  * McGaugh & Bothun (1994) McGaugh, S. S., & Bothun, G. D. 1994, AJ, 107, 530
  * McQuinn et al. (2012) McQuinn, K. B. W., Skillman, E. D., Dalcanton, J. J., et al. 2012, ApJ, 759, 77
  * Mendes de Oliveira et al. (2004) Mendes de Oliveira, C., Cypriano, E. S., Sodré, L., & Balkowski, C. 2004, ApJ, 605, L17
  * Merritt et al. (2014) Merritt, A., van Dokkum, P., & Abraham, R. 2014, ApJ, 787, L37
  * Merritt et al. (2016) Merritt, A., van Dokkum, P., Danieli, S., et al. 2016, ApJ, 833, 168
  * Meurer et al. (2006) Meurer, G. R., Hanish, D. J., Ferguson, H. C., et al. 2006, ApJS, 165, 307
  * Meyer et al. (2004) Meyer, M. J., Zwaan, M. A., Webster, R. L., et al. 2004, MNRAS, 350, 1195
  * Mihos et al. (2018) Mihos, J. C., Durrell, P. R., Feldmeier, J. J., Harding, P., & Watkins, A. E. 2018, ApJ, 862, 99
  * Mihos et al. (2013) Mihos, J. C., Harding, P., Spengler, C. E., Rudick, C. S., & Feldmeier, J. J. 2013, ApJ, 762, 82
  * Mihos et al. (2012) Mihos, J. C., Keating, K. M., Holley-Bockelmann, K., Pisano, D. J., & Kassim, N. E. 2012, ApJ, 761, 186
  * Mould et al. (2000) Mould, J. R., Huchra, J. P., Freedman, W. L., et al. 2000, ApJ, 529, 786
  * Nassau (1945) Nassau, J. J. 1945, ApJ, 101, 275
  * Osterbrock (1989) Osterbrock, D. E. 1989, Astrophysics of gaseous nebulae and active galactic nuclei (University Science Books)
  * Pascale et al. (2021) Pascale, R., Bellazzini, M., Tosi, M., et al. 2021, MNRAS, 501, 2091
  * Pickles (1998) Pickles, A. J. 1998, PASP, 110, 863
  * Pogson (1856) Pogson, N. 1856, MNRAS, 17, 12
  * Román & Trujillo (2017) Román, J., & Trujillo, I. 2017, MNRAS, 468, 4039
  * Rownd et al. (1994) Rownd, B. K., Dickey, J. M., & Helou, G. 1994, AJ, 108, 1638
  * Rudolph et al. (1996) Rudolph, A. L., Brand, J., de Geus, E. J., & Wouterloot, J. G. A. 1996, ApJ, 458, 653
  * Ryan-Weber et al. (2004) Ryan-Weber, E. V., Meurer, G. R., Freeman, K. C., et al. 2004, AJ, 127, 1431
  * Sakai et al. (2002) Sakai, S., Kennicutt, R. C., van der Hulst, J. M., & Moss, C. 2002, ApJ, 578, 842
  * Sandage & Binggeli (1984) Sandage, A., & Binggeli, B. 1984, AJ, 89, 919
  * Schechter (1976) Schechter, P. 1976, ApJ, 203, 297
  * Schmidt (1959) Schmidt, M. 1959, ApJ, 129, 243
  * Schombert et al. (2011) Schombert, J., Maciel, T., & McGaugh, S. S. 2011, AdAst, 2011, 143698
  * Schombert et al. (1990) Schombert, J. M., Wallin, J. F., & Struck-Marcell, C. 1990, AJ, 99, 497
  * Simon (2019) Simon, J. D. 2019, ARA&A, 57, 375
  * Slater et al. (2009) Slater, C. T., Harding, P., & Mihos, J. C. 2009, PASP, 121, 1267
  * Stepanian (2005) Stepanian, J. A. 2005, RMxAA, 41, 155
  * Strömgren (1939) Strömgren, B. 1939, ApJ, 89, 529
  * Sulentic et al. (2001) Sulentic, J. W., Rosado, M., Dultzin-Hacyan, D., et al. 2001, AJ, 122, 2993
  * Taylor et al. (2005) Taylor, V. A., Jansen, R. A., Windhorst, R. A., Odewahn, S. C., & Hibbard, J. E. 2005, ApJ, 630, 784
  * Thilker et al. (2007) Thilker, D. A., Bianchi, L., Meurer, G. R., et al. 2007, ApJS, 173, 538
  * Thuan & Martin (1981) Thuan, T. X., & Martin, G. E. 1981, ApJ, 247, 823
  * Torres-Flores et al. (2009) Torres-Flores, S., Mendes de Oliveira, C., de Mello, D. F., et al. 2009, A&A, 507, 723
  * van der Hulst et al. (1993) van der Hulst, J. M., Skillman, E. D., Smith, T. R., et al. 1993, AJ, 106, 548
  * van der Hulst & Sancisi (1988) van der Hulst, T., & Sancisi, R. 1988, AJ, 95, 1354
  * van Dokkum et al. (2014) van Dokkum, P. G., Abraham, R., & Merritt, A. 2014, ApJ, 782, L24
  * van Dokkum et al. (2015) van Dokkum, P. G., Abraham, R., Merritt, A., et al. 2015, ApJ, 798, L45
  * Vílchez-Gómez (1999) Vílchez-Gómez, R. 1999, in The Low Surface Brightness Universe, ed. J. I. Davies, C. Impey, & S. Phillipps, Vol. 170, ASP, 349
  * Waller et al. (1997) Waller, W. H., Bohlin, R. C., Cornett, R. H., et al. 1997, ApJ, 481, 169
  * Walter et al. (2006) Walter, F., Martin, C. L., & Ott, J. 2006, AJ, 132, 2289
  * Watkins et al. (2017) Watkins, A. E., Mihos, J. C., & Harding, P. 2017, ApJ, 851, 51
  * Weisz et al. (2011a) Weisz, D. R., Dolphin, A. E., Dalcanton, J. J., et al. 2011a, ApJ, 743, 8
  * Weisz et al. (2011b) Weisz, D. R., Dalcanton, J. J., Williams, B. F., et al. 2011b, ApJ, 739, 5
  * Werk et al. (2010) Werk, J. K., Putman, M. E., Meurer, G. R., et al. 2010, AJ, 139, 279
  * Xu et al. (2003) Xu, C. K., Lu, N., Condon, J. J., Dopita, M. A., & Tuffs, R. J. 2003, ApJ, 595, 665
  * Xu et al. (1999) Xu, C. K., Sulentic, J. W., & Tuffs, R. 1999, ApJ, 512, 178
  * Xu et al. (2005) Xu, C. K., Iglesias-Paramo, J., Burgarella, D., et al. 2005, ApJ, 619, L95