# Young, blue, and isolated stellar systems in the Virgo Cluster. II. A new
class of stellar system

Michael G. Jones Steward Observatory, University of Arizona, 933 North Cherry
Avenue, Rm. N204, Tucson, AZ 85721-0065, USA  David J. Sand Steward
Observatory, University of Arizona, 933 North Cherry Avenue, Rm. N204, Tucson,
AZ 85721-0065, USA  Michele Bellazzini INAF – Osservatorio di Astrofisica e
Scienza dello Spazio di Bologna, Via Gobetti 93/3, 40129 Bologna, Italy
Kristine Spekkens Department of Physics and Space Science, Royal Military
College of Canada P.O. Box 17000, Station Forces Kingston, ON K7K 7B4, Canada
Department of Physics, Engineering Physics and Astronomy, Queen’s University,
Kingston, ON K7L 3N6, Canada  Ananthan Karunakaran Instituto de Astrofísica de
Andalucía (CSIC), Glorieta de la Astronomía, 18008 Granada, Spain  Department
of Physics, Engineering Physics and Astronomy, Queen’s University, Kingston,
ON K7L 3N6, Canada  Elizabeth A. K. Adams ASTRON, the Netherlands Institute
for Radio Astronomy, Oude Hoogeveesedijk 4,7991 PD Dwingeloo, The Netherlands
Kapteyn Astronomical Institute, PO Box 800, 9700 AV Groningen, The Netherlands
Giuseppina Battaglia Instituto de Astrofísica de Canarias, Vía Láctea s/n
38205 La Laguna, Spain  Universidad de La Laguna, Avda. Astrofísico Fco.
Sánchez , La Laguna, Tenerife E-38205, Spain  Giacomo Beccari European
Southern Observatory, Karl-Schwarzschild-Strasse 2, D-85748 Garching bei
München, Germany  Paul Bennet Space Telescope Science Institute, 3700 San
Martin Drive, Baltimore, MD 21218, USA  John M. Cannon Department of Physics &
Astronomy, Macalester College, 1600 Grand Avenue, Saint Paul, MN 55105, USA
Giovanni Cresci INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5,
50125 Firenze, Italy  Denija Crnojević University of Tampa, 401 West Kennedy
Boulevard, Tampa, FL 33606, USA  Nelson Caldwell Center for Astrophysics,
Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138, USA  Jackson
Fuson Department of Physics & Astronomy, Macalester College, 1600 Grand
Avenue, Saint Paul, MN 55105, USA  Puragra Guhathakurta UCO/Lick Observatory,
University of California Santa Cruz, 1156 High Street, Santa Cruz, CA 95064,
USA  Martha P. Haynes Cornell Center for Astrophysics and Planetary Science,
Space Sciences Building, Cornell University, Ithaca, NY 14853, USA  John L.
Inoue Department of Physics & Astronomy, Macalester College, 1600 Grand
Avenue, Saint Paul, MN 55105, USA  Laura Magrini INAF - Osservatorio
Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy  Ricardo R.
Muñoz Departamento de Astronomía, Universidad de Chile, Camino El Observatorio
1515, Las Condes, Santiago  Burçin Mutlu-Pakdil Kavli Institute for
Cosmological Physics, University of Chicago, Chicago, IL 60637, USA
Department of Astronomy and Astrophysics, University of Chicago, Chicago IL
60637, USA  Anil Seth Department of Physics & Astronomy, University of Utah,
Salt Lake City, UT, 84112, USA  Jay Strader Center for Data Intensive and Time
Domain Astronomy, Department of Physics and Astronomy, Michigan State
University, East Lansing, MI 48824, USA  Elisa Toloba Department of Physics,
University of the Pacific, 3601 Pacific Avenue, Stockton, CA 95211, USA
Dennis Zaritsky Steward Observatory, University of Arizona, 933 North Cherry
Avenue, Rm. N204, Tucson, AZ 85721-0065, USA

###### Abstract

We discuss five blue stellar systems in the direction of the Virgo cluster,
analogous to the enigmatic object SECCO 1 (AGC 226067). These objects were
identified based on their optical and UV morphology and followed up with H i
observations with the VLA (and GBT), MUSE/VLT optical spectroscopy, and HST
imaging. These new data indicate that one system is a distant group of
galaxies. The remaining four are extremely low mass ($M_{\ast}\sim 10^{5}$
M⊙), are dominated by young, blue stars, have highly irregular and clumpy
morphologies, are only a few kpc across, yet host an abundance of metal-rich,
$12+\log(\mathrm{O/H})>8.2$, H ii regions. These high metallicities indicate
that these stellar systems formed from gas stripped from much more massive
galaxies. Despite the young age of their stellar populations, only one system
is detected in H i, while the remaining three have minimal (if any) gas
reservoirs. Furthermore, two systems are surprisingly isolated and have no
plausible parent galaxy within $\sim$30′ ($\sim$140 kpc). Although tidal
stripping cannot be conclusively excluded as the formation mechanism of these
objects, ram pressure stripping more naturally explains their properties, in
particular their isolation, owing to the higher velocities, relative to the
parent system, that can be achieved. Therefore, we posit that most of these
systems formed from ram pressure stripped gas removed from new infalling
cluster members, and survived in the intracluster medium long enough to become
separated from their parent galaxies by hundreds of kiloparsecs, and that they
thus represent a new type of stellar system.

Low surface brightness galaxies (940); Dwarf galaxies (416); Galaxy
interactions (600); Tidal tails (1701); Ram pressure stripped tails (2126); HI
line (690); Virgo cluster (1772)

††journal: AAS Journals††facilities: Blanco, GALEX, GBT, HST (ACS), ROSAT,
VLA, VLT:Yepun (MUSE)††software: DOLPHOT (Dolphin, 2000), CASA (McMullin et
al., 2007), astropy (Astropy Collaboration et al., 2013, 2018), APLpy
(Robitaille & Bressert, 2012; Robitaille, 2019), Photutils (Bradley et al.,
2020), reproject (Robitaille et al., 2020), acstools (Lim et al., 2020), DS9
(Joye & Mandel, 2003), dustmaps (Green, 2018), Astroalign (Beroiz et al.,
2020), Sextractor (Bertin & Arnouts, 1996), Aladin (Bonnarel et al., 2000;
Boch & Fernique, 2014)

##  1 Introduction

Exceptionally high gas-to-stellar mass ratio systems are of particular
interest in extragalactic astronomy as they represent one extreme of galaxy
formation, namely some of the lowest mass objects that succeed in forming any
stars. Blind radio surveys of neutral hydrogen (H i) have uncovered a plethora
of gas-rich systems that have few, or perhaps no, stars (Saul et al., 2012;
Adams et al., 2013; Taylor et al., 2013; Cannon et al., 2015). However,
distinguishing those which may be genuine extreme low-mass dwarf galaxies from
other classes of objects (Cannon et al., 2015), such as tidal debris and high
velocity clouds (Adams et al., 2016), is a challenging process owing to the
faintness of any associated stellar counterpart to these objects (e.g. Janesh
et al., 2019), as well as confusion with foreground Milky Way H i emission,
which often dominates the velocity range where candidates are expected to be
detectable.

However, attempts to distinguish these objects have led to surprising
discoveries, most notably SECCO 1 (also called AGC 226067; Adams et al., 2013;
Bellazzini et al., 2015; Sand et al., 2015; Adams et al., 2015; Beccari et
al., 2017a; Sand et al., 2017; Bellazzini et al., 2018) and AGC 226178 (Cannon
et al., 2015; Junais et al., 2021; Jones et al., 2022). These are two young,
blue, extremely low-mass ($M_{\ast}\sim 10^{5}$ M⊙), gas-rich, metal-rich,
actively star-forming stellar systems in the Virgo cluster. AGC 226178 has a
gas-to-stellar mass ratio $(1.4M_{\mathrm{HI}}/\mathrm{M_{\ast}})\sim 1000$,
while SECCO 1 has a ratio of $\sim$150.111Here a factor of 1.4 is used to
account for helium in the gas mass. The properties of both systems imply that
they formed via in situ star formation (SF) in gaseous debris stripped from a
much larger object. In the case of AGC 226178, the likely parent object has
been identified as the nearby galaxy VCC 2034, to which it is connected via a
tenuous, low column density, 70 kpc-long H i bridge (Jones et al., 2022).
However, it is unclear whether this gas was stripped by a high-speed tidal
encounter, or by ram pressure from the intracluster medium (ICM). In the case
of SECCO 1, despite it being relatively close to the Virgo cluster center, it
is still sufficiently isolated that its origin is uncertain, and multiple
possible parent objects have been suggested (Sand et al., 2017; Bellazzini et
al., 2018).

As alluded to by Sand et al. (2017) and Jones et al. (2022), these two objects
are not unique but instead appear to be part of a larger population of similar
objects in Virgo. SECCO 1 and AGC 226178 were originally identified through
their H i line emission, thereby guaranteeing gas richness. However, with the
latest and deepest wide field imaging surveys it is possible to visually
identify objects in the Virgo cluster with similar optical and UV properties,
though not necessarily equivalently gas-rich.

In this work, we present comprehensive observations of a sample of isolated,
blue stellar systems in the Virgo cluster as part of a campaign to understand
their physical properties and origins. These additional candidate objects,
along with AGC226067/SECCO1 and AGC226178, were followed up with Hubble Space
Telescope (HST) F606W and F814W imaging with the Advanced Camera for Surveys
(ACS), and H i observations with the Jansky Very Large Array (VLA) and the
Green Bank Telescope (GBT). Additional observations with the MUSE (Multi Unit
Spectroscopic Explorer) integral field spectrograph on the VLT (Very Large
Telescope), are presented in a companion paper, Bellazzini et al. (2022),
hereafter Paper I.

The sample identification is described in §2 and their follow-up observations
in §3. The results, H i and stellar masses, star formation rates (SFRs), and
metallicity measurements are presented in §4. In §5 we search for potential
points of origin of these objects and we discuss potential formation scenarios
in §6. Finally, in §7 & §8 we discuss the fate of these objects and future
directions of investigation, before drawing our conclusions in §9.

We adopt 16.5 Mpc (Mei et al., 2007) as the distance to the Virgo cluster
throughout.

##  2 Target identification

We performed a visual search for isolated, blue stellar systems, similar in
optical appearance to SECCO 1, using the $\sim$100 deg2 of NGVS (Next
Generation Virgo cluster Survey, Ferrarese et al., 2012) $ugi$ imaging of the
Virgo cluster, along with GALEX (Galaxy Evolution Explorer, Martin et al.,
2005) UV imaging when available. Characteristic systems display an over-
density of compact blue sources with strong associated UV emission. They also
lack a diffuse red component typical of Virgo dwarf galaxies, even when they
have ongoing SF. Partial results from this search were presented in Sand et
al. (2017). In total, five isolated, blue stellar system candidates (or BCs),
which we number 1-5, were identified. All five were followed-up with
observations with HST, the VLA, and MUSE/VLT. The coordinates of these five
targets are listed in Table 1, and their locations relative to the Virgo
cluster are shown in Figure 1.

The object we refer to as BC3 is an independent re-identification (based on
optical appearance) of the H i-selected object AGC 226178 from the ALFALFA
survey (Haynes et al., 2011; Cannon et al., 2015). This object has already
been studied in detail (Cannon et al., 2015; Junais et al., 2021; Jones et
al., 2022) and is the BC most similar to SECCO 1.

As discussed in the remainder of this paper, we are now confident that four of
the five BCs are genuine SECCO 1 analogs.

Table 1: BC coordinates and H i velocities

Object | R.A. | Dec. | $v_{\mathrm{HI}}/\mathrm{km\,s^{-1}}$
---|---|---|---
BC1 | 12:39:02.0 | +12:12:16.7 |
BC2‡ | 12:44:27.9 | +12:37:13.4 |
BC3 | 12:46:42.5 | +10:22:04.8 | 1581
BC4 | 12:26:25.7 | +14:23:12.2 |
BC5 | 12:26:30.9 | +15:10:26.2 |
SECCO1 | 12:21:53.9 | +13:27:37.0 | $-142^{\dagger}$

††footnotetext: Columns: (1) object name; (2 & 3) coordinates (J2000) of the
main body of each object; (4) heliocentric velocity of H i emission (Haynes et
al., 2011). †Value for the main body from Adams et al. (2015). ‡BC2 is a
spurious object (§4.1).

##  3 Observations and reduction

After the initial identification of our target BCs using NGVS and GALEX we had
little information about their properties except that they were similar to
SECCO 1 in optical appearance (extremely blue, faint and clumpy) and that
their UV emission indicated some recent or ongoing SF. We therefore pursued a
three-pronged observational strategy to uncover their nature: 1) HST imaging
to better understand their detailed morphology and stellar populations; 2)
Observations with MUSE/VLT to measure their redshifts via the H$\alpha$ line
and obtain metallicity measurements; 3) VLA D-array and GBT observations to
search for any associated H i line emission and quantify their neutral gas
content.

###  3.1 HST observations

Each of the five candidates was observed with ACS in the F606W and F814W
filters as part of program 15183 (PI: D. Sand). Each target was observed for a
total of 2120 s and 2180 s in the two filters respectively, except BC4, which
was observed for 2000 s in each filter. DOLPHOT’s (Dolphin, 2000, 2016) ACS
module was used to align the exposures and perform point source photometry of
the resolved stellar population. The dust maps of Schlegel et al. (1998) and
$R_{\mathrm{F606W}}$ and $R_{\mathrm{F814W}}$ values of Schlafly & Finkbeiner
(2011) were used to correct for Galactic extinction at the position of each
source. Stars were selected from the resulting DOLPHOT catalog following a
similar approach to Jones et al. (2022). Briefly, we select all point-like
(type 1 and 2) objects with no photometry flags from the DOLPHOT source
catalog. We removed sources with greater than 1 mag of crowding (combined,
from the two filters). Finally, the combined (in quadrature) absolute
sharpness value was enforced to be below $\sqrt{0.075}$ and a roundness
threshold of less than 1 (in both filters) was set. Completeness limits were
also estimated as in Jones et al. (2022), based on artificial stars added
evenly over both images in each field. The measured 90% completeness limits
were fit with the combination of a horizontal line and a one-sided parabola
(e.g. Figure 2, bottom panels), and the 50% limits were fit with straight
lines.

In addition to the point source photometry in §4.6 we also perform aperture
photometry on the combined, drizzled images in each band to measure the
integrated magnitudes and colors of the systems. This was performed using the
Astropy package Photutils (Bradley et al., 2020) and a combination of
manually-constructed circular and elliptical apertures enclosing the various
clumps of each source. In each case the sky background was subtracted based on
the median value within an annulus (circular or elliptical) surrounding the
aperture.

###  3.2 MUSE/VLT observations

To robustly identify H ii regions, obtain optical redshifts and basic
kinematics, and measure metallicities, we observed all BCs with MUSE/VLT
(Bacon et al., 2014). These observations were carried out as part of program
0101.B-0376A (P.I: R. Muñoz). They covered the spectral range 4650-9300 Å and
a $\simeq 1.0\arcmin\times 1.0\arcmin$ field centered on each target. These
observations are discussed in detail in Paper I, and here we present an
outline of the data reduction process.

The reduction and analysis of these data followed Beccari et al. (2017a). The
individual dithered exposures were calibrated separately and then combined
into a single stacked data cube for each target. H$\alpha$ (and integrated
light) peaks at least 3$\sigma$ above the background were identified using
Sextractor (Bertin & Arnouts, 1996). The flux of each of these detected
sources was measured using a 1.5″ (radius) aperture and a 1D spectrum (with a
step size of 1.25 Å) of each source was produced.

Redshifts were measured for all detected H$\alpha$ clumps, and line fluxes for
H$\beta$, [Nii], and [Oiii] were measured wherever possible (Tables 2 & 3 of
Paper I, ). In Section 4.2 we summarize the findings of these measurements and
their implications for the origins of BCs.

###  3.3 GALEX data

We searched for archival NUV and FUV data from GALEX at the location of each
BC (and SECCO 1). Most of the BCs are within the footprint of the GALEX
Ultraviolet Virgo Cluster Survey (GUViCS, Boselli et al., 2011), however, this
is not always the deepest data available. For BCs 2, 3, 4 and SECCO 1 we use
tiles from GUViCS (typically $\sim$1.6 ks in both bands), but no FUV data are
available for either BC4 or SECCO 1. For BC1 we use tiles “Virgo_Epoque_MOS05”
($\sim$16 ks) and “NGA_Virgo_MOS04” ($\sim$1.6 ks) for NUV and FUV,
respectively. For BC5 we use tiles “NGA_NGC4421” ($\sim$2 ks) and
“GI1_079012_Group5” ($\sim$1.6 ks).

In §4.5 we perform aperture photometry on these GALEX tiles and estimate the
SFR in each object. The flux within each aperture was measured from the
corresponding background subtracted GALEX tile. Uncertainties were estimated
by placing 10000 circular apertures (equal in area to the target apertures)
randomly across the GALEX tile after masking the brightest 1% of pixels.
Magnitudes were calculated following the conversions of Morrissey et al.
(2007) and extinction corrections used $R_{\mathrm{NUV}}=8.20$ and
$R_{\mathrm{FUV}}=8.24$ (Wyder et al., 2007). Finally, these magnitude
measurements were converted to SFRs following Iglesias-Páramo et al. (2006),
using 4.74 as the bolometric solar absolute magnitude.

###  3.4 VLA observations

BC3 was observed previously as part of the ALFALFA “Almost Dark” galaxies
sample (VLA program 13A-028, PI: J. Cannon; Cannon et al., 2015). These data
were obtained in D-configuration, have a channel width of 7.81 kHz ($\sim$1.65
km s-1), a total bandwidth of 8 MHz, and a total on-source integration time of
approximately 1.6 h. These data were re-reduced by Jones et al. (2022) using
standard reduction methods in the Common Astronomy Software Applications
package (CASA, McMullin et al., 2007). The final imaging used Brigg’s
robust=0.5 weighting to provide a compromise between sensitivity and angular
resolution for the detected H i emission. The channels were averaged and re-
binned to a velocity resolution of 5 km s-1.

The remaining 4 candidates were observed in the VLA program 18A-185 (PI: K.
Spekkens). Each target was observed on-source for approximately 1.5 h in
D-configuration. The initial observations of both BC1 and BC2 suffered from
severe interference and were subsequently re-observed, greatly improving the
data quality. As the redshift of the objects were not known prior to the
observations, we used a 32 MHz bandwidth (from 1394.416 to 1426.416 MHz, or
approximately $-$1250 to 5500 km s-1) to search for any H i emission
associated with the optical candidates. This range was split up into 3072
channels of 10.42 kHz ($\sim$2.2 km s-1), which during the data reduction was
averaged over 4 channels resulting in a velocities resolution of 8.8 km s-1.

Initially the entire bandwidth of the data were reduced to search for H i
emission. However, after the redshifts for all candidates were obtained from
MUSE spectroscopy with the VLT, a narrow sub-band was re-reduced (spanning
$\sim$1000 km s-1), allowing for improved local continuum subtraction. The
reduction was performed with a Python and CASA-based pipeline that will be
presented in full in Jones et al. (in prep.). The most severe interference was
flagged manually and the tfcrop flagging algorithm was run in addition. For
BC1 we also used rflag, after initial calibrations, as there were no bright
lines which might be mistaken for interference (other than Milky Way
emission). Imaging used Brigg’s robust=2 weighting in order to maximize our
detection capabilities. Refer to Table 2 for details of the beam sizes and rms
noise for each observation.

Table 2: VLA data summary

Object | Beam size | $\sigma_{\mathrm{rms}}/\mathrm{mJy\,beam^{-1}}$ |  $\Delta v_{\mathrm{chan}}/$km s-1
---|---|---|---
BC1 | 60″$\times$51″ | 1.9 | 8.8
BC2 | 63″$\times$55″ | 1.1 | 8.8
BC3 | 56″$\times$45″ | 1.2 | 5
BC4 | 65″$\times$54″ | 0.9 | 8.8
BC5 | 65″$\times$54″ | 1.0 | 8.8

††footnotetext: Columns: (1) object name; (2) synthesized beam size; (3) rms
noise; (4) velocity resolution.

###  3.5 GBT observations

The large surface area and low system temperature of the GBT allow it to
obtain much deeper H i spectra than the VLA, providing a more stringent
constraint on any neutral gas content. However, after the redshifts of the
candidates were known (from their H$\alpha$ emission) it was determined that
only BC1 was suitable for single dish follow-up as BC4 and BC5 would be
confused with Milky Way emission, BC3 had already been strongly detected with
the VLA (Cannon et al., 2015), and the HST imaging of BC2 indicated that it
was a background galaxy group (§4.1). A director’s discretionary time proposal
(21A-433, PI: M. Jones) was submitted to the GBT and BC1 was observed for a
total of 3 h using $\mathrm{ON-OFF}$ position switching. The data were reduced
using standard GBT IDL procedures. The resulting spectrum has an rms noise of
0.25 mJy (within $\pm 300$ km s-1 of the redshift of BC1) after smoothing to a
velocity resolution of 30 km s-1.

##  4 Results

Figure 1: Locations of BCs (and SECCO 1) in the direction of Virgo overlaid on
a ROSAT mosaic of hard (0.4-2.4 keV) X-ray emission (Brown et al., 2021).
Virgo members and possible members (from the Extended Virgo Cluster Catalog;
Kim et al., 2014) are plotted as faint black, unfilled circles. The area of
each circle is proportional to the total $r$-band flux of the galaxy it
represents. The BCs are shown with blue symbols (see legend) and SECCO 1 is
shown as a purple cross. The symbol for BC2 is unfilled as this object is
spurious (see §4.1). The approximate virial radius (taken to be 1.7 Mpc,
Kashibadze et al., 2020) of the cluster is shown by a large dashed black
circle.

Figure 2: Top-left: False color HST F606W+F814W image of BC1. The dashed green
circle shows the region used to construct the CMD. At the distance of the
Virgo cluster (16.5 Mpc) 20″ is 1.6 kpc. Top-right: GALEX NUV+FUV image
showing the same field. Bottom-left: CMD of the point sources within the
aperture shown. The dashed lines indicates the 90% completeness limit and the
dotted line the 50% limit. The errorbars indicate the typical uncertainties
(from artificial star tests) in the F814W magnitude and F606W-F814W color, as
a function of F814W magnitude. Bottom-right: The CMD of a background region of
the HST image away from bright sources. The aperture used was equal in area to
the target aperture.

Figure 3: Top-left: False color HST F606W+F814W image of BC2. The dashed green
circle shows the region used to construct the CMD. Unlike the other BCs, this
HST images appears to indicate that this is a background galaxy group. Top-
right: GALEX NUV+FUV image showing the same field. There is only very weak NUV
emission associated with BC2. Bottom: CMD within the aperture shown (left) and
a blank field aperture (right). The dotted and dashed lines, and error bars,
are the same as described in Figure 2, bottom panels. This CMD appears to be
consistent with background, supporting the conclusion that this is a spurious
blue stellar system candidate.

Figure 4: Top-left: False color HST F606W+F814W image of BC3. The dashed green
ellipse and circles show the regions used to construct the CMD. Top-right:
GALEX NUV+FUV image showing the same field. Bottom: CMD within the apertures
shown (left) and a blank field aperture (right). See Figure 2 caption for
further details.

Figure 5: Top: False color HST F606W+F814W image of BC4. The dashed green
ellipses and circles show the regions used to construct the CMD. Middle: GALEX
NUV image showing the same field. Bottom: CMD within the apertures shown
(left) and a blank field aperture (right). See Figure 2 caption for further
details.

Figure 6: Top-left: False color HST F606W+F814W image of BC5. The dashed green
ellipse and circle show the regions used to construct the CMD. The component
BC5c was identified via H$\alpha$ emission (Paper I) to be at the same
velocity as the main body, but may only be a single cluster of stars. Top-
right: GALEX NUV+FUV image showing the same field. Bottom: CMD within the
apertures shown (left) and a blank field aperture (right). See Figure 2
caption for further details. Figure 7: H i spectra of BC1 from the GBT (top)
and the VLA (bottom). The VLA spectrum was extracted using an aperture equal
in area to the synthesized beam. The vertical dashed lines correspond to the
H$\alpha$ velocity measurement from MUSE. No significant signal is detected in
either spectrum. Figure 8: H i spectra of BC3 from ALFALFA and the VLA. The
ALFALFA spectrum is the public spectrum from Haynes et al. (2018) and the VLA
spectrum was created using the extended source mask of Jones et al. (2022).
The vertical dashed line corresponds to the H$\alpha$ velocity measurement
from MUSE. BC3 is detected at high signal-to-noise ratio in both spectra and
both agree with the H$\alpha$ velocity. However, the VLA measures a somewhat
lower flux, with most of the missing emission lying on the approaching side of
the line profile. This likely indicates the presence of extended emission
below the surface brightness limit of the VLA observations (Cannon et al.,
2015; Jones et al., 2022).

Figure 9: H i spectra in the directions of BC4 (top) and BC5 (bottom)
extracted from the VLA H i data cubes within an aperture equal in area to the
synthesized beam. The vertical dashed lines correspond to the H$\alpha$
velocity measurements from MUSE. Neither shows a significant H i line signal,
although they could be contaminated with Milky Way H i emission. The apparent
peak coincident with the H$\alpha$ velocity of BC5 is below 3$\sigma$ and is
likely a noise spike.

In this section we present the results of our multi-wavelength investigation
of BCs, providing a description of their morphology, colors, redshifts,
stellar masses, metallicities, and gas content. These physical properties will
then be used as the basis for a search for candidate parent objects in the
following section and a discussion of potential formation pathways in §6.

We include SECCO 1 in this sample throughout and either use quantities
measured in previous work or (re)measure them as needed (e.g. to provide
equivalent values across the whole sample).

###  4.1 Morphology and location

The HST images of the BCs are shown in Figures 2, 3, 4, 5, 6, while that for
SECCO 1 can be found in Sand et al. (2017). In all cases, except BC2, the BCs
appears to be very blue, highly irregular, and frequently broken up into
multiple components. Their stellar populations also appear to be partially
resolved with some individual stars discernible. These are almost exclusively
blue and likely only represent the youngest, brightest stars, not the
underlying stellar population. However, their extremely blue appearance
(discussed further in §4.4 and §4.6) also suggests that any redder, underlying
population is likely minimal.

The largest single component of any BC is BC3a (Figure 4, top-left), which is
approximately 30″ across its major axis (2.4 kpc at the distance of Virgo).
However, what constitutes a single component is quite subjective, for example,
BC4a, b, and c could justifiably be considered as a single object (Paper I).
The smallest components (e.g. BC3c, BC5b, BC5c) are less than 5″ ($\sim$400
pc) across and may only consist of a single cluster of stars.

The six components of BC4 are spread over $\sim$1.5′ ($\sim$7 kpc) and may
indicate that this is either a very young collection of objects formed in a
gas-rich stream, or a somewhat older object that has become gravitationally
unbound. Even for the other BCs, which are mostly defined by one or two
components, their highly irregular and clumpy structure points to them being
extremely low-mass and potentially unbound. Although the individual components
of the BCs were identified visually, based primarily on the HST images, nearly
all of these clumps have corresponding UV and (usually) H$\alpha$ emission
(Paper I), which indicate ongoing SF. In the case of the latter, the
components are all kinematically associated (see §4.2).

In the case of BC2, the HST image (Figure 3, top-left) is quite distinct from
the other BCs and indicates that this is a spurious candidate. It appears to
be a distant background group of galaxies rather than a nearby young object.
Unlike the other BCs there is also minimal UV emission (particularly FUV)
associated with this candidate (Figure 3, top-right), and it was undetected in
H$\alpha$ by MUSE (§4.2). Furthermore, almost no stars were identified in its
CMD (Figure 3, bottom-left), which is consistent with background (Figure 3,
bottom-right). Henceforth, we will not regard BC2 as a genuine blue stellar
system and statements regarding the global properties of BCs should be assumed
to include SECCO 1, but not BC2.

Figure 1 shows the locations of the BCs on the sky in relation to Virgo
cluster galaxies and the cluster virial radius. BC2 is shown as an unfilled
symbol. All of the BCs are within the virial radius of the cluster. However,
none are in the very cluster center, within $\sim$2∘ ($\sim$575 kpc) of M 87
(the central galaxy in the Virgo cluster, Figure 1). BC1 is the closest, with
a projected separation of approximately 600 kpc. This may indicate that the
parent objects of BCs are recent additions to the cluster.

###  4.2 H$\alpha$ velocities and metallicities

Table 3: Metallicities of BCs

Object | $v_{\mathrm{H}\alpha}/\mathrm{km\,s^{-1}}$ | $N_{\mathrm{H}\alpha}$ | $N_{\mathrm{O/H}}$ | $\langle 12+\log\mathrm{O/H}\rangle$
---|---|---|---|---
BC1 | $1117\pm 6$ | 18 | 2 | $8.35\pm 0.15$
BC3 | $1584\pm 4$ | 15 | 5 | $8.29\pm 0.17$
BC4 | $-60\pm 19$ | 16 | 6 | $8.73\pm 0.15$
BC5 | $-74\pm 5$ | 4 | 2 | $8.70\pm 0.14$
SECCO1† | $-153.2\pm 1.4$ | 33 | 9 | $8.38\pm 0.11$

††footnotetext: H$\alpha$ redshift and metallicity measurements from Paper I.
Columns: (1) object name; (2) mean velocity (and standard deviation) of
H$\alpha$ clumps detected with MUSE; (3) number of clumps detected in
H$\alpha$; (4) number of clumps detected in H$\alpha$, H$\beta$, [Nii], and
[Oiii] (suitable for deriving an O/H estimate); (5) mean oxygen abundance and
uncertainties (standard deviation of clumps and scatter in O3N2 calibration,
0.14 dex). †Values from Beccari et al. (2017a).

MUSE detected H$\alpha$ emission in all BCs, identifying between 4 and 18
distinct clumps of emission in each object (Paper I). The mean velocity (and
standard deviation) of these clumps in each source is shown in Table 3. Only
BC3 (and SECCO 1) has a prior velocity from an H i detection (Table 1), which
matches closely with the H$\alpha$ velocity for that object. All the objects
have velocities which are consistent with Virgo cluster membership
($-500<cz_{\odot}/\mathrm{km\,s^{-1}}<3000$, e.g. Mei et al., 2007), and all
are (projected) within the virial radius of the cluster (Figure 1). We note
that although BC4, BC5, and SECCO 1 all have negative radial velocities, they
are in the vicinity of M 86 ($cz_{\odot}=-224$ km s-1), a region of the Virgo
cluster where negative radial velocities are common.

As described in Paper I, the average oxygen abundance of each BC was estimated
based on N2 and O3N2 (following Pettini & Pagel, 2004), which were corrected
for extinction based on the relative strengths of H$\alpha$ and H$\beta$. The
resulting metallicity estimates are shown in Table 3. All the BCs are
extremely high metallicity given their very low stellar masses (§4.4), which
suggests that they formed from gas pre-enriched in more massive objects. Of
particular note are BC4 and 5, both of which are found to be marginally super-
solar in metallicity ($12+\log(\mathrm{O/H})_{\odot}=8.69$, Asplund et al.,
2009). The details of the kinematics and metallicity spreads of the clumps
within the BCs are discussed in Paper I.

###  4.3 H i mass & limits

Table 4: H i masses of BCs

Object | $M_{\mathrm{HI}}/\mathrm{M_{\odot}}$ | Telescope
---|---|---
BC1 | $<1.6\times 10^{6}$ | GBT
BC3 | $4.0\times 10^{7}$ | Arecibo†
BC4 | $<2.9\times 10^{6}$ | VLA
BC5 | $<3.2\times 10^{6}$ | VLA
SECCO1 | $1.5\times 10^{7}$ | Arecibo‡

††footnotetext: Columns: (1) object name; (2) H i mass or 3$\sigma$ upper
limit; (3) telescope for the stated value. †Haynes et al. (2011). ‡Adams et
al. (2015).

SECCO 1 is the prototype BC, first detected via its H i emission (Adams et
al., 2013), having a total H i mass of $1.5\times 10^{7}$ M⊙ (Adams et al.,
2015). The low resolution of H i observations means that the main and
secondary body of SECCO 1 appear as one source in H i. However, Adams et al.
(2015) also identified an additional H i-only component slightly to the north
as well as another potential optical component, also to the north, but not
coincident with any H i. Owing to the similar optical/UV appearance of BCs 1-5
it was anticipated that they would also be H i-rich, which motivated our VLA
follow-up program.

In Figures 7, 8, and 9 we present the VLA (and GBT) H i spectra of BCs 1, 3,
4, and 5 (the spectrum of BC2 is discussed in Appendix A). The VLA H i spectra
were extracted from the data cubes using an aperture equal to the synethesized
beam size, centered on the location of the main body of each BC (Table 1). In
addition to these spectra the data cubes were visually inspected channel by
channel and SoFiA was run to search for significant emission features that
might be extended spatially or spectrally.

Like SECCO 1, BC3 was known a priori to contain a significant H i reservoir as
it was originally identified in the ALFALFA survey (Haynes et al., 2011).
However, among BCs 1-5 this is the only object that was detected in our VLA
observations. Based on the VLA spectrum (Figure 8, extracted using the SoFiA
source mask), and an assumed distance of 16.5 Mpc, BC3 has an H i mass of
$\log M_{\mathrm{HI}}/\mathrm{M_{\odot}}=7.3$. This value is 0.3 dex lower
than that measured by ALFALFA (Haynes et al., 2011) suggesting that the VLA
has not recovered all the extended flux (this was also noted by Cannon et al.,
2015; Jones et al., 2022). Jones et al. (2022) show that when viewed in the
ALFALFA data cube (which has better column density sensitivity for extended
emission than the VLA observations) the H i emission coincident with BC3
appears to connect to the galaxy VCC 2034, approximately 70 kpc to the SW.
This galaxy is almost certainly the source of the gas that formed BC3
(discussed further in §5, and Jones et al., 2022).

If the other BCs had comparable H i masses to BC3 and SECCO 1, then they would
be detected with the VLA observations, but none were (Figures 7 & 9). The
slight caveat is that, because of their low radial velocities, BC4 and BC5
might be blended with MW H i emission. The spectrum of BC5 (Figure 9, bottom)
also appears to have a peak coincident with the H$\alpha$ velocity of BC5.
However, this peak is below 3$\sigma$ and extremely narrow, and is likely a
noise spike.

All the BCs that are undetected in H i have optical redshift measurements from
MUSE H$\alpha$ observations (BC3 and SECCO 1 do also) and the available H i
data can therefore be used with confidence to set upper limits on their H i
masses. For BC4 and BC5 the deepest data are those from the VLA, which have
rms noise values of 0.9 and 1.0 mJy/beam (at 8.8 km s-1 resolution),
respectively, at the velocities of the H$\alpha$ emission. Assuming that any H
i emission would fit within one synthesized beam (Table 2) and would have a
velocity width of 30 km s-1, then these equate to 3$\sigma$ upper limits of
$\log M_{\mathrm{HI}}/\mathrm{M_{\odot}}<6.46$ and 6.51, respectively,
assuming a fiducial distance of 16.5 Mpc in both cases. For BC1 the GBT follow
up spectrum is by far the more sensitive. With an rms of 0.28 mJy (at 30 km
s-1 resolution) this gives the 3$\sigma$ upper limit as $\log
M_{\mathrm{HI}}/\mathrm{M_{\odot}}<6.2$, again assuming a fiducial distance of
16.5 Mpc. These limits are listed in Table 4.

###  4.4 Stellar populations

Table 5: Magnitudes and stellar mass estimates

Object | F814W | F606W-F814W | $M_{\ast}/\mathrm{M_{\odot}}$
---|---|---|---
BC1 | $20.29\pm 0.38$ | $0.08\pm 0.41$ | $\sim 5\times 10^{4}$
BC3 | $20.23\pm 0.15$ | $-0.23\pm 0.17$ | $\sim 5\times 10^{4}$
BC4 | $19.86\pm 0.26$ | $-0.26\pm 0.29$ | $\sim 1\times 10^{5}$
BC5 | $20.56\pm 0.10$ | $0.06\pm 0.12$ | $\sim 5\times 10^{4}$
SECCO1 | $20.39\pm 0.41$ | $-0.23\pm 0.46$ | $\sim 4\times 10^{4}$

††footnotetext: Columns: (1) object name; (2) F814W magnitude (extinction
corrected); (3) F606W-F814W color (extinction corrected); (4) stellar mass
estimate (§4.4).

Figure 10: Reproduced CMDs of BCs 1, 3, 4, and 5, with PARSEC isochrones for
different stellar population ages overlaid, assuming a distance of 16.5 Mpc to
all objects. The isochrones for BC1 and 3 use a metallicity of $[M/H]=-0.35$,
which is approximately the value for both objects. For BC4 and 5 the value is
$[M/H]=0.05$. In the latter case it is possible that the two objects are
associated, while the similar metallicities of BC1 and 3 are likely by chance.
In the leftmost panel we also plot an isochrone indicating where an old (10
Gyr) RGB population would reside in these diagrams, below the completeness
limit in the lower right corner and barely visible.

The HST (and GALEX) images and associated CMDs for all BCs are shown in
Figures 2, 3, 4, 5, and 6. The blue optical colors and UV emission indicate
that the BCs have predominantly young, blue stellar populations. Furthermore,
the detection of H$\alpha$ in all BCs indicates that the youngest stars must
be $\leq$10 Myr old.

As discussed by Jones et al. (2022), the CMD of BC3 (Figure 4, bottom) is most
similar to that of SECCO 1 (Sand et al., 2017), apparently made up almost
entirely of blue main sequence (MS) and helium burning stars (F814W $\gtrsim$
24.5, F606W-F814W $\lesssim$ 0) and red helium burning (RHeB) stars (23.5
$\lesssim$ F814W $\lesssim$ 26.5 and F606W-F814W $\gtrsim$ 0.6 mag), with
almost no candidates for red giant branch (RGB) stars, highlighting the young
age of the population.

BC1’s CMD (Figure 2, bottom) is similar, but the brightest RHeB stars are more
numerous and fainter than in BC3. These slight differences likely indicate
that BC1 is somewhat older than BC3 (as RHeB peak brightness is a function of
age, e.g. McQuinn et al., 2011), which would be consistent with its non-
detection in H i, if sufficient time has passed for its neutral gas to have
been evaporated or stripped.

The CMDs of BC4 and BC5 (Figures 5 & 6, bottom panels) are again similar, but
the RHeB stars are even fainter and continue to the completeness limit. This
likely indicates that BC4 and BC5 are the oldest objects in the sample. The
color spread between the blue and RHeB stars is also wider for BC4 and BC5
than for any of the other BCs.

In Figure 10 we overplot PARSEC isochrones (PAdova and TRieste Stellar
Evolution Code, Bressan et al., 2012) on the CMDs of each object for a variety
of stellar population ages. As pointed out by Jones et al. (2022) the faintest
RHeB stars in BC3 appear consistent with the 50 Myr isochrone, likely
indicating that the stellar population in this object cannot be much older
than 50 Myr. In the case of the other BCs, as mentioned above, their CMDs
imply that their oldest stars are somewhat older (although they must still
have formed young stars within the past 10 Myr, as they contain H ii regions),
but the proximity of the RHeB stars to the completeness limit prevents them
from being used to estimate ages. The isochrones also explain the different
color gap between the bluest and reddest stars in the CMDs of BC4 and BC5
versus BC1 and BC3. This spread is approximately reproduced in the isochrones
and is a function of the higher metallicity of these two objects, which
despite their feeble appearance is marginally super-solar (Table 3).

The CMD of SECCO 1 was presented and discussed in Sand et al. (2017) and
Bellazzini et al. (2018). The general appearance is similar to the other BCs.
The spread between the reddest and bluest stars is most similar to BC1 and
BC3, again a reflection of the similar metallicities of these objects. Sand et
al. (2017) also simulate a mock stellar population and argue that SECCO 1 must
be younger than $\sim$50 Myr based on the luminosity of the RHeB stars. This
is roughly consistent with our estimate of 60 Myr in §4.6 based on the
integrated F814W magnitude and SFR of SECCO 1.

If we compare the CMDs of the BCs to the low-mass, gas-rich dwarf Leo P
(McQuinn et al., 2015b), then we see a striking difference. In addition to the
young blue stars in Leo P there is also a clear, well-populated RGB at a
similar magnitude that is entirely absent from the BCs CMDs. This clear RGB is
the result of the old underlying population in Leo P, but for extremely young
stellar populations (which BCs appear to be) no RGB population exists.
Furthermore, any RGB stars would be significantly less luminous than the young
stars that dominate the CMDs of the BCs. However, the proximity of Leo P
($D=1.6$ Mpc) means that the depth of its CMD is a mis-match for those of BCs,
making it an unfair comparison, despite it being one of the most similar
objects known in terms of SFR and stellar mass (but notably not metallicity).

A fairer comparison can be made by considering a blue, irregular dwarf at the
distance of Virgo, in this case VCC 1816 (KDG 177, $M_{V}=-15.2$), which has
similar depth HST observations as the BCs. The CMD of this galaxy (Figure 4 of
Karachentsev et al., 2014) shows both a blue population (at
F606W-F814W$\sim$0) and a red population (at F606W-F814W$\sim$1), similar to
BCs. The former is likely made up of blue helium burning stars and young MS
stars, as in the BCs, while the latter is likely made up of a combination of
asymptotic giant branch and RHeB stars. The number of stars in the CMD
increases towards fainter F814W magnitudes (near F606W-F814W$\sim$1) probably
indicating the presence of a well-populated RGB near the completeness limit,
which is lacking in the BC CMDs. A similar lack of evidence for any RGB was
noted for SECCO 1 by Sand et al. (2017) and Bellazzini et al. (2018), but in
comparison to a red dwarf spheroidal in Virgo, rather than a star-forming
dwarf more in line with the appearance of BCs.

At the distance of the Virgo cluster the tip of the red giant branch (TRGB) is
expected to be at F814W $\sim$ 27 mag (e.g. Jiang et al., 2019), which would
be borderline detectable with our HST observations. However, at high
metallicities the TRGB becomes less defined and RGB stars become redder, both
of which would impede the detectability of an RGB in our observations (Figure
10, leftmost panel). Thus, it is not possible to conclusively rule out there
being an RGB based on the CMDs of BCs. Despite this we still view the
existence of an underlying old population as extremely unlikely in these
objects. They are extremely blue, to the point where stellar population models
struggle to reproduce their colors, even when assuming very young ages (§4.6).
In addition, BCs were specifically selected (§2) to be lacking any visible
diffuse red component in the deep NGVS images. Together these points make it
highly unlikely that there could be any significant underlying old population
of stars, even though the CMDs themselves are insufficiently deep to reach any
potential RGB.

Overall the CMDs of the BCs can be characterized as having a population of
stars made up exclusively of young blue main sequence and (blue and red)
helium burning stars, with no evidence of a RGB. The luminosities of the RHeB
stars suggest that the youngest BCs (BC3 & SECCO 1) are around 50 Myr old.
Finally, the remarkably high metallicities measured with MUSE (§4.2) appear to
be consistent with the color difference between the reddest and bluest helium
burning branch stars in the CMDs.

###  4.5 Star formation rates

SFRs (Table 6) were estimated for each candidate by measuring the NUV and FUV
fluxes within the same apertures used to produce their CMDs (§4.4), as
described in §3.3. An additional uncertainty of 15% was added to the error
budget as this is the stated accuracy of the conversion in Iglesias-Páramo et
al. (2006).

The SFRs of all BCs fall in the range
$-3.5<\log\mathrm{SFR/M_{\odot}\,yr^{-1}}<-3$ and are generally quite
consistent between NUV and FUV (where both images are available), likely
indicating that their SFRs have not varied strongly over the past $\sim$100
Myr (or that they are younger than this). Although matching NUV and FUV SFR
estimates could be the result of a bursty SF histories over the past $\sim$100
Myr, with an average rate that equals that of the past $\sim$10 Myr, it seems
highly unlikely that this could be the case for all BCs, and a constant SFR is
a more natural explanation for this finding. The UV-based SFRs are also
roughly consistent with the SFRs estimated from the integrated H$\alpha$
fluxes (Paper I), with the slight exception of BC1 (for which the SFR may be
beginning to decline), again supporting the assertion that the SFRs appear to
have been relatively constant in the recent past. We note that had we adopted
a different conversion scheme for our UV-based SFR estimates (e.g. McQuinn et
al., 2015a) then our SFRFUV values could be up to 0.6 dex higher. However,
given the general consistency between the H$\alpha$ and UV-based SFR
estimates, the conversion scheme we originally selected appears appropriate
for these objects.

This range of SFRs is similar to the faintest dwarf irregular galaxies in the
Local Volume (Lee et al., 2009). However, the extremely low stellar masses of
BCs (§4.6) make it difficult to directly compare to equivalent star-forming
dwarf galaxies, as almost none are known at these masses. For example, even
Leo P (Giovanelli et al., 2013) is almost an order of magnitude higher stellar
mass than most BCs, but its SFR is around an order of magnitude lower
($\log\mathrm{SFR/M_{\odot}\,yr^{-1}}=-4.4$ McQuinn et al., 2015b). Leo T
(Irwin et al., 2007) is of comparable stellar mass to BCs ($M_{\ast}=1.4\times
10^{5}$ M⊙, Weisz et al., 2014), but is apparently no longer forming stars, or
is between episodes (Kennicutt et al., 2008). If we put the SFRs of BCs in
terms of their specific SFRs (sSFR) then they fall in the range
$-8.2<\log(\mathrm{SFR}/M_{\ast})/\mathrm{yr^{-1}}<-7.7$, which would place
them significantly higher than average, but within the scatter, of sSFR for
low-mass, gas-rich, field galaxies (Huang et al., 2012; James et al., 2015).

Table 6: UV fluxes and SFR estimates

Object |  | $\mathrm{SNR}_{\mathrm{NUV}}$ | NUV flux | $\log\frac{\mathrm{SFR_{NUV}}}{\mathrm{M_{\odot}\,yr^{-1}}}$ | $\mathrm{SNR}_{\mathrm{FUV}}$ | FUV flux | $\log\frac{\mathrm{SFR_{FUV}}}{\mathrm{M_{\odot}\,yr^{-1}}}$ | $\log\frac{\mathrm{SFR_{H\alpha}}}{\mathrm{M_{\odot}\,yr^{-1}}}$ | $\log\frac{M_{\mathrm{HI}}/\mathrm{SFR}}{\mathrm{yr}}$
---|---|---|---|---|---|---|---|---|---
BC1 |  | 10.1 | $9.10\pm 0.90$ | $-3.25\pm 0.08$ | 13.7 | $1.98\pm 0.14$ | $-3.42\pm 0.07$ |  $-$3.9 | $<9.5$
BC3 |  |  | $17.4\pm 0.7$ | $-3.03\pm 0.07$ |  | $4.02\pm 0.09$ | $-3.18\pm 0.07$ |  $-$3.1 | 10.3
| a | 23.8 | $14.1\pm 0.6$ | $-3.13\pm 0.07$ | 43.9 | $3.36\pm 0.08$ | $-3.26\pm 0.07$ |  |
| b | 9.3 | $2.78\pm 0.30$ | $-3.83\pm 0.08$ | 13.6 | $0.57\pm 0.04$ | $-4.03\pm 0.07$ |  |
| c | 4.6 | $0.57\pm 0.12$ | $-4.52\pm 0.11$ | 5.6 | $0.10\pm 0.02$ | $-4.77\pm 0.10$ |  |
BC4 |  |  | $13.1\pm 0.7$ | $-3.10\pm 0.07$ |  |  |  |  $-$3.2 | $<9.6$
| a | 16.5 | $4.07\pm 0.25$ | $-3.60\pm 0.07$ |  |  |  |  |
| b | 4.7 | $1.04\pm 0.22$ | $-4.20\pm 0.11$ |  |  |  |  |
| c | 11.8 | $2.37\pm 0.20$ | $-3.84\pm 0.08$ |  |  |  |  |
| d | 6.1 | $1.49\pm 0.24$ | $-4.04\pm 0.10$ |  |  |  |  |
| e | 7.4 | $3.56\pm 0.48$ | $-3.66\pm 0.09$ |  |  |  |  |
| f | 4.1 | $0.55\pm 0.13$ | $-4.47\pm 0.13$ |  |  |  |  |
BC5 |  |  | $6.13\pm 0.31$ | $-3.48\pm 0.07$ |  | $1.20\pm 0.05$ | $-3.69\pm 0.07$ |  $-$3.8 | $<10.0$
| a | 20.1 | $5.57\pm 0.28$ | $-3.52\pm 0.07$ | 27.4 | $1.13\pm 0.04$ | $-3.72\pm 0.07$ |  |
| b | 2.8 | $0.38\pm 0.14$ | $-4.68\pm 0.16$ | 2.8 | $0.06\pm 0.02$ | $-4.96\pm 0.16$ |  |
| c | 3.7 | $0.18\pm 0.05$ | $-4.99\pm 0.12$ | 1.3 | $0.01\pm 0.01$ | $-5.67\pm 0.33$ |  |
SECCO1 |  |  | $10.4\pm 0.8$ | $-3.14\pm 0.07$ |  |  |  |  $-$3.2† | $10.3$
| MB | 11.1 | $6.72\pm 0.06$ | $-3.33\pm 0.08$ |  |  |  |  |
| SB | 7.6 | $3.63\pm 0.48$ | $-3.60\pm 0.09$ |  |  |  |  |

††footnotetext: Columns: (1) object name and sub-component (where relevant);
(2) SNR of NUV emission (see Section 4.5 for details); (3) NUV flux in units
of $10^{-17}\;\mathrm{erg\,s^{-1}\,cm^{-2}\,\AA^{-1}}$; (4) NUV-based SFR
estimate; (5) SNR of FUV emission; (6) FUV flux in units of
$10^{-16}\;\mathrm{erg\,s^{-1}\,cm^{-2}\,\AA^{-1}}$; (7) FUV-based SFR
estimate; (8) H$\alpha$ SFR estimates from the integrated H$\alpha$ flux of
each object in MUSE (Paper I) following the conversion of Kennicutt (1998);
(9) gas consumption timescale using the larger of the NUV and FUV SFR
estimates (we note that this quantity is distance independent). For
uniformity, all objects are assumed to be at 16.5 Mpc. †H$\alpha$-based SFR
estimate from Beccari et al. (2017a, b).

###  4.6 Stellar masses

The integrated F606W and F814W magnitudes of the BCs were measured from the
co-added images in each filter. The same apertures indicated in Figures 2, 4,
5, and 6 were used to measure the total magnitude of each source after masking
the few clear background galaxies contained within these apertures. Galactic
extinction corrections were made using the dust maps of Schlegel et al. (1998)
and the reddening $R_{\nu}$ values of Schlafly & Finkbeiner (2011). The final
magnitudes are listed in Table 5. The uncertainties were estimated by placing
10 apertures across the full ACS FoV (avoiding bright stars and background
galaxies) and using the standard deviation of the counts to approximate the
uncertainty in the counts of each BC.

In young stellar populations the emitted light is dominated by the youngest
stars, but the mass is generally dominated by the oldest, most numerous stars.
As BCs are apparently such young objects, the correct mass-to-light ratio to
use is highly uncertain, and would depend strongly on the assumed age of each
object. Thus, widely used mass-to-light ratio prescriptions (e.g. Zibetti et
al., 2009; Taylor et al., 2011) cannot be used with confidence for such a
young, irregular, low-mass, and metal-rich stellar population. We therefore
adopt an unconventional strategy for estimating the stellar masses of the BCs.
If the current SFRs are assumed to be reasonable representations of the SFRs
over the (short) lifetimes of the BCs then the total stellar mass is simply
the age of each object times its SFR. In order to estimate the age we build up
the integrated F814W magnitude of a stellar population forming stars at a
constant rate (in 10 Myr steps), based on the PARSEC (Bressan et al., 2012)
population models. When the artificial F814W magnitude equals the measured
magnitude, we obtain an age estimate for the BC in question (to the nearest 10
Myr).

To estimate the ages, and subsequently the stellar masses (age $\times$ SFR),
we used the NUV SFR measurements (Table 6) for each object, as these are
available for all objects and reflect a slightly longer SF timescale. For BC1,
BC3, and SECCO1 a metallicity of $[M/H]=-0.35$ was used, and $[M/H]=0.05$ for
BC4 and BC5. These values approximately correspond to their observed O/H
values (Table 3). The age estimates222We note that these age estimates should
be treated with caution. Ideally the full SF histories of the BCs would be
calculated, but the currently existing data are inadequate to do this. These
ages represent an approximation to the age of the oldest stellar population in
each BC, based on the assumption of a roughly constant SFR. for BCs 1, 3, 4,
5, and SECCO 1 are 90, 50, 110, 160, and 60 Myr, and the resulting stellar
mass estimates are shown in Table 5.

A significant caveat to this approach is that the PARSEC models are incapable
of correctly reproducing the colors of the BCs (as noted by Sand et al.,
2017). Although this issue is not fully addressed in this work, we chose to
rely on the F814W magnitudes as the discrepancy is assumed to be most severe
for the youngest, bluest stars. Hence the redder band is expected to be
somewhat less impacted. We also note that there are encouraging trends in the
values that we obtained, that at least indicate internal consistency. For
example, BC3 and SECCO 1 are the only BCs detected in H i and we estimate
these are by far the youngest objects — a finding that the CMDs of the BCs
would also seem to support. The estimated ages of BC4 and BC5 are also the
oldest and it seems plausible (§B.2) that the two formed from the same origin.
In addition, Junais et al. (2021) independently estimated the stellar mass of
BC3 by fitting the spectral energy distribution (from photometry in $ugriz$,
H$\alpha$, NUV, and FUV) of each clump with a single stellar population, via a
grid search over metallciity and population age, and found a near identical
value ($\sim 5\times 10^{4}$ M⊙).

##  5 Points of origin

Figure 11: V-band luminosity versus metallicity (relative to solar) for BCs,
Local Group dwarfs (Kirby et al., 2013), Local Volume dwarfs (Berg et al.,
2012), TDGs (Duc & Mirabel, 1998; Weilbacher et al., 2003; Duc et al., 2007;
Croxall et al., 2009; Lee-Waddell et al., 2018), and extremely metal-poor
galaxies (XMPs, Skillman et al., 2013; McQuinn et al., 2015b; Hirschauer et
al., 2016; Hsyu et al., 2017; Izotov et al., 2019; McQuinn et al., 2020).
Metallicity is measured either from Fe/H or O/H as indicated in the legend.
Both BCs and TDGs sit well above the luminosity–metallicity relation for dwarf
galaxies of equivalent luminosities.

Figure 12: All VCC, EVCC, and ALFALFA neighbors of BC1, 3, 4, and 5 (top-left
to bottom-right) in a 4 sq deg region centered on the BC (blue star in each
panel) and within $\pm$500 km s-1 of their H$\alpha$ velocity (Table 3). The
area of each circular marker corresponds to the apparent magnitude (in
$g$-band) of the galaxy it represents. The color of the markers corresponds to
the galaxy’s $g-i$ color, with the narrow transition from blue to red
occurring at $g-i=0.9$ (green). Objects circled with a dashed blue line were
detected in ALFALFA and thus contain significant quantities of H i (note that
BC3 itself was detected in ALFALFA, but is not circled here). At the distance
of Virgo 30′ corresponds to $\sim$140 kpc. The scale bar in the bottom-right
panel applies to all panels. Figure 13: As for Figure 12, but for SECCO 1.

The observations presented above reveal the surprising result that although
all the BCs are actively forming stars, only BC3 and SECCO 1 have a detectable
quantity of H i. However, the typical values of the SFR estimates are of the
order of $10^{-3}\;\mathrm{M_{\odot}\,yr^{-1}}$ (Table 6), which means that
even below our H i detection limits (§4.3) these objects could still have gas
consumption timescales in excess of 1 Gyr (although such long timescales are
not uncommon for low-mass galaxies). In addition, the high metallicities of
BCs (Table 3) clearly point to them all having formed from pre-enriched gas
that originated in a larger galaxy, as has been shown explicitly to be the
case for BC3 (Jones et al., 2022), where the gas trail can still be traced
back to its parent galaxy.

Figure 11 compares the metallicities of BCs to other objects of comparable
luminosity. As expected, tidal dwarf galaxies (TDGs) are similar to BCs, being
of equivalent metallicity, but typically somewhat higher luminosities. This is
a point to which we will return in the following section (§6), but BCs are
likely too low-mass to be TDGs. On the opposite end of the metallicity
spectrum we compare BCs to a small selection of extremely metal-poor galaxies
(XMPs). These objects can appear superficially similar to BCs. Both are
usually extremely blue, have clumpy morphologies, and the faintest XMPs are
the same luminosity as BCs. However, their metallicities could scarcely be
more different and the two populations are clearly distinct in origin.

If we take the metallicities of the BCs and use them to infer a stellar mass
from the mass–metallicity relation (MZR) then this should provide a reasonable
estimate of the type of galaxies from which they formed. Using the MZR of
Andrews & Martini (2013), the metallicities of the BCs imply that their parent
objects could have stellar masses anywhere in the range $8.3\lesssim\log
M_{\ast}/\mathrm{M_{\odot}}\lesssim 10.1$ (we note that because the MZR is an
asymptotic relation, the lower bound of this range is much better constrained
than the upper bound). This covers a broad range from dwarf galaxies almost to
Milky Way-like galaxies ($\log M_{\ast,\mathrm{MW}}/\mathrm{M_{\odot}}=10.8$,
Licquia & Newman, 2015), but all are massive enough that, unless particularly
low surface brightness (LSB), they should be mostly included in existing
catalogs of Virgo cluster galaxies. Furthermore, it is also reasonable to
assume that the parent objects are gas-bearing (or were in the recent past),
as they must have been able to supply the gas that formed the young stellar
populations of the BCs. The quoted range is for the average metallicity and
does not account for metallicity variations within the parent galaxies, which
could potentially expand the range if the material that formed a BC originated
from a region that strongly deviated from the average metallicity.

###  5.1 Search for points of origin

We performed a detailed search considering all known Virgo members in the
vicinity (and at a similar velocity to) each BC, paying particular attention
to gas-bearing galaxies detected in ALFALFA (Haynes et al., 2018). Even though
most BCs are undetected in H i, they appear to have formed from stripped gas
and must have contained gas in the recent past as they have all formed stars
recently. Thus, nearby, gas-rich galaxies are good candidate progenitor
systems. The galaxies neighboring each BC are shown in Figures 12 & 13. Here
we present the conclusions of our search, but the full details can be found in
Appendix B.

Within the entire 4 deg2 region shown around BC1 in Figure 12 (top-left),
there is only one galaxy that contains H i gas and is sufficiently massive to
have formed a BC, NGC 4579. This galaxy is approximately 140 kpc (30′) to the
SW of BC1 and separated from it in velocity by $\sim$400 km s-1. Thus, BC1
would have needed a large ejection velocity ($>$500 km s-1 in total) for NGC
4579 to be its parent object. Furthermore, other than a slightly H i-deficient
disk333Chung et al. (2009) define the H i-deficiency of a galaxy’s disk as the
logarithmic decrement between the observed and expected mean H i surface
density within the optical disk, where for the latter they use the average
value for isolated galaxies from Haynes & Giovanelli (1984)., NGC 4579 shows
little sign of recent disturbance in either its H i or CO morphology and
kinematics (Chung et al., 2009; Brown et al., 2021). Finally, NGC 4579 appears
to be too metal-rich ($12+\log(\mathrm{O/H})=8.87\pm 0.05$, De Vis et al.,
2019) to match the metallicity of BC1 ($\langle
12+\log(\mathrm{O/H})\rangle=8.35\pm 0.15$). Thus, NGC 4579 does not seem to
be a viable candidate point of origin for BC1, and the genuine point of origin
must presumably be beyond the region shown in Figure 12 ($>280$ kpc away), but
we are unable to identify any strong candidates this far away.

BC3 (also called AGC 226178) was discussed in detail by Jones et al. (2022).
This is an extremely complicated field with multiple foreground systems
projected on it. BC3’s nearest apparent neighbor is NGVS 3543 (also called AGC
229166), which Junais et al. (2021) argued was a LSB galaxy at the same
distance as BC3. However, based on the CMDs produced from HST imaging, Jones
et al. (2022) demonstrated that NGVS 3543 is a foreground object at $\sim$10
Mpc, while BC3 is consistent with being in Virgo at 16.5 Mpc. H i observations
with the VLA (Cannon et al., 2015) and Arecibo (Giovanelli et al., 2005;
Haynes et al., 2011; Minchin et al., 2019) indicate a possible bridge between
BC3 and a pair of galaxies, VCC 2034 and 2037. However, the closer (in
projection) of these, VCC 2037, is actually another foreground object at
approximately 10 Mpc (Karachentsev et al., 2014). Thus, VCC 2034
($cz_{\odot}=1507$ km s-1), $\sim$70 kpc to the SW, is the likely source of
BC3’s H i gas. However, Jones et al. (2022) were unable to determine whether
ram pressure or tidal stripping was responsible for removing the gas from VCC
2034.

BC4 and BC5 likely formed from the same parent object as they are only
separated by 45′ on the sky, are at almost the same velocity (Table 3), have
nearly identical metallicity measurements (Table 3), and have similar age
estimates (§4.6). NGC 4419 is in fairly close proximity to both BCs and based
on its estimated stellar mass and metallicity, it is likely a close match for
the metallicity of these BCs. In addition, there is strong evidence in both H
i (Chung et al., 2009) and CO (Brown et al., 2021) that this galaxy is being
ram pressure stripped. However, as ram pressure tails only extend in one
general direction (in the wake of a galaxy’s motion through the ICM) and BC4
is to the south of NGC 4419 and BC5 is to the north, it is extremely unlikely
that this is the point of origin of these BCs. The extension of the molecular
gas distribution of NGC 4419 is roughly towards the south (Brown et al.,
2021), in the direction of BC4, but away from BC5. If NGC 4419 were
simultaneously undergoing ram pressure and tidal stripping then it could
plausibly have formed both BCs, but there is no evidence of this in the
optical, H i, or CO images.

There are are a few other gas-bearing galaxies within 1 deg of either BC4 or
BC5, but these were all discounted due to a mismatch in properties or because
of evidence showing ram pressure stripping in the wrong direction. Upon
searching further afield we immediately identified UGC 7695 (VCC 1450, IC
3476) as a strong candidate. This galaxy is a well-studied example of ram
pressure stripping in action (Boselli et al., 2021b), and has a prominent bow-
shaped wake extending in the approximate direction of BC4 and BC5. Existing
measurements of the metallicity of UGC 7695 (Hughes et al., 2013; Boselli et
al., 2021b) also approximately match those of BC4 and BC5, making this a
promising candidate for their point of origin. As the separation between the
BCs and UGC 7695 is approximately 450 kpc in projection they would presumably
require a very large (perhaps over 1000 km s-1) ejection velocity, depending
on when the stripping episode began.

Finally, we consider SECCO 1. Figure 13 shows the neighbors of SECCO 1 within
a 4 deg2 field and $\pm$500 km s-1, and demonstrates the extraordinary
isolation of this system given that it is within the virial radius of a
cluster. The potential points of origin for SECCO 1 have already been
discussed extensively by previous works (Adams et al., 2015; Sand et al.,
2017; Bellazzini et al., 2018) and we will only review these briefly here.

If formed by a stripping event then the most likely point of origin is either
the M 86 subgroup of Virgo, about 350 kpc to the SE, which exhibits an
enormous complex of stripped gas visible in X-rays and H$\alpha$ (Sand et al.,
2017, and references therein), or the group of dwarf galaxies $\sim$200 kpc to
the NW (Bellazzini et al., 2018). The proximity of VCC 322, 334, and 319
(compared to the M 86 sub-group) might favor this possibility. However, as we
have discussed above, in some cases the separation between parent and BC may
be quite large. What is a stronger argument is that the metallicities of VCC
322 and 334 are a close match to that of SECCO 1 (Bellazzini et al., 2018).
VCC 322 also has a stellar tail that extends in the general direction of SECCO
1. Although these galaxies are less massive than some of the others
considered, we note that the apparent parent object of BC3 is also a dwarf
galaxy and only a few times more massive than BC3 itself.

However, the complex of stripped gas (e.g. Boselli et al., 2018) in the M 86
sub-group is also a good candidate point of origin, for example, if NGC 4438
(VCC 1043, beyond the FoV shown in Figure 13) fell towards this sub-group via
the location of SECCO 1. In this case, a combination of ram pressure and tidal
forces could be responsible for SECCO 1 and the complex of stripped gas. We
also noted the blue dwarf irregular IC 3355 near NGC 4438 (in the approximate
direction of SECCO 1). However, the lower metallicity of this object
($12+\log(\mathrm{O/H})\approx 8.0$, De Vis et al., 2019) suggests that it did
not form from stripped gas.

###  5.2 Other origin scenarios

In the above discussion we considered that BCs were likely formed from a gas-
bearing galaxy sufficiently massive to be included in existing catalogues of
Virgo cluster galaxies. However, there are a few other scenarios that we
briefly consider here.

Junais et al. (2021) suggested that BC3 might have formed from gas stripped
from a LSB galaxy. Although this scenario is ruled out for BC3 itself (as the
LSB galaxy in question is actually a foreground object, Jones et al., 2022) it
is possible that LSB galaxies have been missed in our search above, as they
are frequently absent from established catalogs of cluster members. In general
this mechanism would imply that the LSB galaxy being stripped would be
relatively close by to the BC, as it would presumably have a smaller gas
reservoir (that would evaporate more rapidly when stripped) than a larger
galaxy. Therefore, even though LSB galaxies can be challenging to detect, it
seems unlikely that a close neighbor would have been overlooked in multiple
cases, and we do not consider this a likely formation pathway, but note that
it is difficult to entirely exclude.

An additional scenario that we considered is the possibility that the BCs
could be dark objects that contained neutral gas for an extended period, but
formed essentially no stars until very recently (c.f. Kent et al., 2009;
Minchin et al., 2019). This scenario is highly unlikely for two main reasons.
Firstly, the search for bona fide dark galaxies that cannot be explained as
tidal or spurious objects has turned up few convincing results to date (e.g.
Taylor et al., 2013; Cannon et al., 2015), calling into question whether this
scenario is valid. Secondly, the high metallicity of the BCs indicates that
there have been multiple prior SF episodes that have enriched their gas, thus
ruling out that they could be primordial dark objects (c.f. Corbelli et al.,
2021).

##  6 Formation mechanism

As shown in Figure 11, the universally high metallicities of BCs (in relation
to their luminosities or stellar masses) means that the only plausible
mechanism for their formation is that they formed from material stripped from
a larger galaxy. Their metallicities are a full order of magnitude higher than
those of galaxies of the same V-band luminosity, and owing to their extremely
young stellar populations, this discrepancy would be even larger if the
samples were compared in terms of their stellar masses. Figure 11 also
indicates that BCs are of slightly lower luminosity than TDGs, but of similar
metallicity. Their stellar and H i masses indicate that BCs are considerably
less massive than long-lived TDGs.

Despite the strong evidence that BCs formed from stripped material, it is
unclear whether they formed through tidal or ram pressure stripping. As is the
case for BC3, even when the H i connection to the parent galaxy is still
detectable (Jones et al., 2022), it may not be possible to distinguish between
ram pressure or tidal forces as the dominant mechanism stripping the gas.
Indeed it is possible that both are valid mechanisms.

Regardless of the mechanism by which gas is stripped to form BCs, the parent
objects must be new cluster members. Oman & Hudson (2016) and Oman et al.
(2021) simulated the stripping and quenching of galaxies falling into clusters
and found that essentially all new members are stripped of their gas and
quenched during their first orbit, usually around pericenter passage (see also
Cortese et al., 2021; Boselli et al., 2021a, for reviews). Thus to have
sufficient gas to form a BC, the parent galaxy must likely be on its first
infall into the cluster. Although we only have a sample of five objects, this
would also appear to agree with their spatial distribution which is inside the
virial radius, where significant stripping is expected, but avoids the cluster
center.

In the remainder of this section we discuss the evidence for and against tidal
and ram pressure formation scenarios, compare BCs to other classes of objects
known to form from stripped gas, and give an overview of related simulation
results.

###  6.1 Comparison to TDGs and the need for ram pressure stripping

The typical masses of long-lived TDGs are expected to be over $10^{8}$ M⊙
(Bournaud & Duc, 2006), as below this mass they generally cannot resist the
tidal field of their parent galaxies for long enough to escape as bound
objects. This threshold mass is considerably larger than any of the BCs,
disfavoring a tidal formation pathway, as the most massive BCs (BC3 and SECCO
1) are a few times 107 M⊙. However, we note that if a lower mass TDG were to
be ejected at particularly high speed it may be able to survive, as it would
more rapidly escape the tidal field of its parent galaxy. We also note that
the simulations of Bournaud & Duc (2006) assume that the parent object of TDGs
are loosely MW-like, however, VCC 2034 (the apparent parent object of BC3,
Jones et al., 2022) has a stellar mass of only $10^{8.2}$ M⊙. If BCs are tidal
in origin then perhaps they formed from lower mass progenitors and are
correspondingly lower mass than typical long-lived TDGs. However, such a
mechanism could presumably only apply to those BCs (BC1, BC3, and SECCO 1)
with slightly lower metallicities that correspond to similarly low-mass
progenitors (via the MZR), unless the more metal-rich BCs (BC4 and BC5) formed
from recently enriched gas that was stripped before it had sufficient time to
mix with the rest of the interstellar medium in the parent galaxy.

TDGs are typically ejected at around the circular velocity of the galaxy they
originate from (Bournaud & Duc, 2006). For a relatively massive galaxy this
might mean an ejection velocity of $\sim$300 km s-1. If this were
preferentially aligned along the direction perpendicular to the line-of-sight,
it would still take a TDG $\sim$1 Gyr to traverse 300 kpc in projection. Thus
the isolation of BCs, coupled with their very young stellar populations, is
difficult to explain via a tidal formation mechanism. In contrast, in the case
of ram pressure stripping the velocity of the galaxies relative to the cluster
can exceed 1000 km s-1, and galaxies with the largest tails are generally
found to be traveling at highest speeds (Jaffé et al., 2018). The fact that
BCs have been identified in a cluster also points to ram pressure stripping as
the most likely formation pathway. All gas-rich galaxies falling with
sufficient velocity into a cluster are expected to undergo some degree of ram
pressure stripping, and while tidal interactions are certainly commonplace in
clusters, these most frequently take the form of brief, high speed encounters
(e.g. “galaxy harassment”), which are less likely to strip large quantities of
gas (e.g. Smith et al., 2010) than the strong, drawn out interactions in
galaxy groups (where TDGs are typically found).

However, although the relative velocity between an infalling galaxy and the
ICM can easily exceed 1000 km s-1, this does not necessarily translate into an
equivalent velocity for the stripped gas, as once stripped it does not
immediately become stationary relative to the ICM. The ram pressure stripping
simulations of Kapferer et al. (2009) consider gas-rich galaxies falling at
1000 km s-1 relative to an ICM of varying densities (from $1\times 10^{-28}$
to $5\times 10^{-27}\;\mathrm{g\,cm^{-3}}$). They show that after 500 Myr of
stripping the length of the plume of stripped gas in the wake of the parent
galaxy is strongly dependent on the density of the surrounding ICM (e.g. their
Figure 20). In this case the most distant gas clouds (for
$\rho_{\mathrm{ICM}}\geq 1\times 10^{-27}\;\mathrm{g\,cm^{-3}}$) are $\sim$400
kpc from their parent galaxy, indicating that their average relative velocity
over the past 500 Myr has been $\sim$800 km s-1. This is somewhat slower than
the velocity of the parent galaxy relative to the ICM (and would be
considerably slower still for lower ICM densities), but is still several times
greater than the relative velocities typically expected for TDGs. For
comparison, the electron number density of the ICM in Virgo (Nulsen &
Bohringer, 1995) exceeds $10^{-2}\;\mathrm{cm^{-3}}$ ($\sim 2\times
10^{-26}\;\mathrm{g\,cm^{-3}}$) near M 87 and at a distance of 230 kpc has
decreased to $6\times 10^{-4}\;\mathrm{cm^{-3}}$ ($\sim 1\times
10^{-27}\;\mathrm{g\,cm^{-3}}$). Thus ram pressure stripping (in Virgo)
provides a more viable mechanism for rapidly achieving large separations
between stripped material and its parent object. This is especially true
within a few hundred kpc of the cluster center, but could be true almost
anywhere within the cluster should an infalling galaxy collide with a dense
pocket in the ICM (which are known to exist in other clusters, e.g. Morandi &
Cui, 2014; Eckert et al., 2015).

The most similar known objects to BCs are “fireballs” (e.g. Cortese et al.,
2007; Yoshida et al., 2008; Hester et al., 2010), clumps of SF seen in the
wake of galaxies being actively ram pressure stripped. Indeed, as discussed by
Bellazzini et al. (2018), many of the properties of BCs match well with those
of fireballs (e.g. Fumagalli et al., 2011), including their metallicities
(e.g. fireballs in the wake of IC 3418 have $8.22<12+\log(\mathrm{O/H})<8.38$,
Kenney et al., 2014). Other related objects include SF clumps in filamentary
structures in the vicinity of NGC 1275 in the Perseus cluster (Conselice et
al., 2001; Canning et al., 2014) and in stripped material in Stephan’s Quintet
(Mendes de Oliveira et al., 2004).

However, BCs are distinct from fireballs and similar objects, in that they are
remarkably isolated (e.g. Figures 12 & 13). Fireballs are found within a few
10s of kpc of their parent galaxy, where there can be little doubt over their
point of origin, and where they may still eventually fall back onto their
parent galaxy (e.g. Vollmer et al., 2001; Tonnesen & Bryan, 2012a). To form
BCs requires a mechanism which can carry neutral gas several 100s of kpc from
a galaxy within the hostile environment of a cluster.

###  6.2 Properties of ram pressure stripped gas clumps in simulations

Lee et al. (2022) argue that many of the molecular gas clouds seen in the tail
of ram pressure stripped galaxies (e.g. Moretti et al., 2018; Jáchym et al.,
2019) could form in situ, by rapid cooling of warm ionized gas (Tonnesen &
Bryan, 2012b; Moretti et al., 2020, also suggest a similar mechanism). The
metal-rich gas and absence of young stars in close proximity (unlike within
the disks of most gas-rich galaxies) make conditions favorable for radiative
cooling. Furthermore, Müller et al. (2021) argue that magnetic sheathing could
help to protect ram pressure stripped gas from evaporation in a cluster
environment.

In the radiative hydrodynamical simulations of Lee et al. (2022) SF clumps are
seen out to $\sim$100 kpc from the parent galaxy. In their models this SF in
the distant tail occurs $\sim$200 Myr after the initial onset of ram pressure
stripping, suggesting that the very young ages (50-150 Myr) of the stellar
populations of BCs may underestimate how long ago their progenitor gas was
stripped (although large velocities $>$500 km s-1 would likely still be
required to explain their isolation). Finally they note that although bright
H$\alpha$ clumps will only track SF activity, more diffuse H$\alpha$ emission
(fainter than $6\times 10^{38}\;\mathrm{erg\,s^{-1}\,kpc^{-2}}$) is expected
throughout the ram pressure tail due to recombinations in the warm ionized
gas. Therefore, H$\alpha$ observations significantly more sensitive than this
threshold might be capable of robustly identifying the points of origin of
BCs. The nominal 1$\sigma$ surface brightness sensitivity of the VESTIGE
survey is $2\times 10^{-18}\;\mathrm{ergs\,s^{-1}\,cm^{-2}\,arcsec^{-2}}$
(Boselli et al., 2018), which for a distance of 16.5 Mpc equates to $3.4\times
10^{36}\;\mathrm{erg\,s^{-1}\,kpc^{-2}}$. Thus, such features, should they
exist, would be detectable in VESTIGE.

The hydrodynamic simulations of Kapferer et al. (2009) also produce numerous
gas clumps in the wakes of ram pressure stripped galaxies, but out to much
greater distances ($\sim$400 kpc). They find that SF is only induced in these
clumps if the wind speed exceeds 500 km s-1 and that it is enhanced by yet
stronger ram pressure (but note that Tonnesen & Bryan, 2012b, find the
opposite trend). Ram pressure therefore appears to be a promising candidate
for producing clumps of star forming gas far from their parent galaxies, but
do the physical properties of these systems match with those observed in BCs?

In the case of the above mentioned ram pressure simulations the gas clumps in
the wakes of the stripped galaxies are generally presented in terms of gas
density rather than masses of distinct clumps. However, in the simulations of
Tonnesen & Bryan (2021) the masses of such clumps are found to be on the order
of $10^{5}$ M⊙, with the most massive distinct clouds being $\sim$106 M⊙. This
matches quite well with the masses of gas clumps typically found in the
immediate wakes of ram pressure stripped galaxies (e.g. Poggianti et al.,
2019), but is more than an order of magnitude less massive than BC3 and SECCO
1. However, the earlier ram pressure stripping simulations of Kronberger et
al. (2008) do form bound objects, analogous to TDGs, with total masses of
$\sim$107 M⊙, but these simulations are now thought to oversimplify fluid
instabilities (e.g. Sijacki et al., 2012), calling into question the details
of these results.

It terms of metallicity it is generally assumed that, in either the tidal or
ram pressure stripping scenario, the BC formed will exhibit the same
metallicity as its parent galaxy. However, Tonnesen & Bryan (2021) also find
that all their simulated ram pressure stripped clouds rapidly mix with the
ICM. Thus, they predict that the metallicity of ram pressure stripped clouds
should decrease with distance from their parent galaxy. This seems to be
directly contradicted by the high metallicities of BCs, given their relative
isolation and large separations (e.g. $>$300 kpc in some cases) from their
apparent points of origin. We also note that Calura et al. (2020) find that
more massive H i clouds (similar to SECCO 1 and BC3) can survive intact for on
the order of a Gyr, while moving rapidly through the ICM. It may be that the
gas clouds from which BCs form are exceptional objects and not typical of the
underlying population of gas clouds that are stripped in ram pressure events.
For example, these could be some of the most loosely bound gas that is the
first to be stripped, or they could be stripped by a denser clump of the ICM.
This could explain the lack of similar objects in the simulations of Tonnesen
& Bryan (2021).

As a closing remark for this discussion we also note that despite the
apparently simple requirement for ram pressure stripping (i.e. sufficient ram
pressure to overcome the gravitational attraction of the gas disk) and
extensive efforts to simulate this process in increasing detail, there remain
systems that are challenging to explain. In particular, the recent discovery
of an enormous (apparently) ram pressure stripped H i tail in a system outside
of a cluster, where no significant inter-galatic medium could be detected
(Scott et al., 2022), poses difficult questions regarding its origin, and
could even suggest that some BC-like objects might exist outside of clusters.

###  6.3 Summary

In summary, BCs appear to be distinct from both TDGs and fireballs, being too
low-mass to be the former, too high-mass to be the latter, and too isolated
for either. The isolation of some BCs is their property that is hardest to
explain, and would seem to necessitate the large velocities expected in ram
pressure stripping events, but not for strong tidal interactions. We therefore
favor ram pressure stripping as the most likely formation mechanism of BCs. If
ram pressure stripping is confirmed to be the formation pathway then BCs can
be thought of as “ram pressure dwarfs”, analogous to tidal dwarfs, but
unlikely to survive as bound structures on long timescales. Simulations
provide a somewhat conflicting picture of how such objects might form via ram
pressure stripping, however, this may be because BCs represent atypical
objects that, unlike fireballs, are not formed in large numbers during
stripping episodes. Regardless of their formation mechanism the properties of
BCs appear to be distinct from any other stellar systems of which we are
aware.

##  7 Fate and production rate

Based on their morphologies and stellar mass estimates, BCs are unlikely to be
gravitationally bound. In the case of BC3 and SECCO 1, their H i content might
be sufficient for them to remain bound in the short term (e.g. Calura et al.,
2020), but this neutral gas (the majority of their total mass) will eventually
be lost to the ICM. It is challenging to accurately assess the stability of
BCs due to their irregular morphologies and because their velocity dispersions
are not well resolved by the MUSE observations. However, based on the stellar
mass estimates in Table 5 and their apparent sizes, we estimate that if they
are extremely dynamically cold (e.g. $\sigma_{v}<1$ km s-1) then they may be
bound, but for $\sigma_{v}>2$ km s-1 they would certainly be unbound (using
Equation 8 of Calura et al., 2020). The stability of BCs is considered further
in Paper I, however, the most likely scenario is that each BC as a whole is
unbound, but individual component clumps or star clusters may be bound, if
sufficiently dynamically cold. Thus, in the long term BCs will likely disperse
(either as individual stars or star clusters) into the intracluster light.
However, even if BCs were to remain bound, without sustained SF they would
quickly become essentially undetectable. Currently they are only identifiable
at all (in optical/UV) because of their young, blue stars.

In Jones et al. (2022) we argued that as BCs are only expected to be visible
for a short period, they must be continually produced in the cluster. We
estimated that an object such as BC3 might be detectable for at most 500 Myr,
meaning that for five BCs to be visible today, they must be being produced at
a rate on the order of 1 per 100 Myr. However, given that all five of the
known BCs appear to have ages of less than 200 Myr, this might be a more
reasonable estimate, making the production rate closer to 1 per 50 Myr. We
speculate that this might be a common phenomenon with many newly infalling
galaxies producing such objects.

With this in mind we also note that, in hindsight, the metallicities of BCs
are perhaps not surprising. As mentioned in §5, the metallicities of BCs
correspond to a stellar mass range $8.3\lesssim\log
M_{\ast}/\mathrm{M_{\odot}}\lesssim 10.1$. However, this wide range likely
encompasses most galaxies that could possibly form BCs (suggesting it is a
common occurrence). Galaxies significantly less massive than $\log
M_{\ast}/\mathrm{M_{\odot}}=8.3$ would have H i reservoirs scarcely larger
than those of SECCO 1 and BC3, and are thus probably too small to form a BC
themselves. Whereas galaxies significantly more massive than $\log
M_{\ast}/\mathrm{M_{\odot}}=10.1$ are increasingly uncommon and increasingly
likely to be gas-poor.

##  8 Future directions

Although the faintness and peculiar properties of BCs make them challenging
objects to study, we suggest a few directions where progress could likely be
made.

We experienced significant difficulties in attempting to identify the point of
origin of most of the BCs, likely because it has been several hundred Myr
since some of them were first stripped from their parent galaxy. However, if
ram pressure stripping is the formation pathway of these objects then it is
possible that extremely faint H$\alpha$ trails still connect the BCs to their
parent objects (e.g. Lee et al., 2022). A deep H$\alpha$ search around the BCs
(or others identified in the future), is a promising approach to robustly
identifying their parent objects, which in turn will allow for more detailed
study of their formation mechanism. X-ray emission is also frequently found to
accompany H$\alpha$ tails of ram pressure stripped galaxies (Sun et al., 2007,
2021, e.g.), and may represent another means to characterize the properties of
the stripping events that formed the BCs in cases where the parent object can
be identified.

CO observations with the Atacama Large Millimeter/submillimeter Array (ALMA)
have successfully made detections of individual clumps of molecular gas in the
wakes of a ram pressure stripped galaxies (e.g. Jáchym et al., 2019) as well
as individual giant molecular clouds in TDGs (Querejeta et al., 2021). If BCs
still contain significant quantities of molecular gas (as would be expected
based on their recent SF) then they should be readily detectable with ALMA,
especially as their high metallicity measurements imply a favorable CO-to-H2
conversion factor in comparison to other low-mass objects (Bolatto et al.,
2013).

Lee et al. (2022) find that ram pressure stripped gas clouds may travel for
over a hundred Myr before SF occurs in them. Although we are limited by a very
small sample size, all of the BCs with slightly older stellar populations
estimates ($\gtrsim 100$ Myr, rather than $\sim$50 Myr) are undetected in H i.
The long gas consumption (by SF) timescales in Table 6 indicate that the gas
in these systems is not (for the most part) consumed by SF. BCs 1, 4, and 5
must have contained significant cold gas reservoirs within the past 200 Myr to
permit the formation of their observed stellar populations, and they likely
still contain some molecular gas as they have all formed new stars in the past
10 Myr, yet today we find no evidence of any H i content. If BCs have traveled
through the ICM for significantly longer than the current age of their stellar
populations in order to reach their current state of isolation, perhaps this
indicates that it is the SF episode itself that triggers the evaporation of
the neutral gas. For example, it seems plausible that an object like SECCO 1
is essentially an earlier stage of an object like BC4 and BC5. Both are broken
into two main components, but SECCO 1 has yet to lose its gas, and BC4 appears
as though it may be disintegrating (Figure 5). If this were the case then
SECCO 1 would likely be on the verge of losing its H i gas. This disagrees
somewhat with the findings of Calura et al. (2020). However, those authors
note that the details of the SF episode, particularly when it began, are quite
uncertain, and we suggest that this possibility might warrant further
investigation.

Along similar lines, with the HST observations it has only been possible to
characterize the stellar populations of BCs from the stars formed in the past
$\sim$50 Myr. While some BCs may genuinely contain no stars that are older
than this, others might. To detect or rule out this older stellar population,
and therefore to constrain the full SF histories of BCs, is possible with the
James Webb Space Telescope observations (JWST). With a moderate investment of
observing time ($\sim$10 h) JWST is capable of detecting stars several
magnitudes below the TRGB, should an RGB exist, at the distance of the Virgo
cluster. Such observations would not only determine the age of the oldest
stellar component of BCs, but (if RGB stars exist in BCs) would also be
capable of conclusively demonstrating Virgo membership via TRGB distance
measurements.

Finally, we note that if our hypothesis is correct and BCs are commonly
produced when new member galaxies fall into a cluster, then they should exist
in other clusters as well as Virgo. Unfortunately, due to how faint BCs are,
they would be undetectable in any galaxy clusters significantly farther away
than Virgo. We therefore suggest that the Fornax cluster could be a suitable
location to extend the search and would represent an independent environment
where our findings could be cross-checked. The distance modulus for Fornax is
only $\sim$0.5 mag greater than for Virgo, thus the brightest BCs (Table 5)
would likely still be detectable and slightly longer HST observations could
provide similar quality CMDs.

The ongoing MeerKAT Fornax survey (Serra et al., 2016; Kleiner et al., 2021)
aims to map 12 deg2 of the cluster in H i. These observations will be
approximately 5 times deeper than our pointed VLA observations of BCs, and
will have a factor of $\sim$3 times better angular resolution. Thus this
survey will be ideal for identifying “dark”, dense H i clouds analogous to
those in Adams et al. (2013) and Cannon et al. (2015) that led to the
discovery of BCs. Furthermore, the improved column density sensitivity will
exceed that of ALFALFA and will be readily capable of detecting residual H i
streams that might still connect young BCs to their parent objects, as is the
case for BC3 (Jones et al., 2022). The Fornax cluster is also the target of
both the Next Generation Fornax Survey (Muñoz et al., 2015), $ugi$ imaging
with the Dark Energy Camera on the Blanco telescope, and the Fornax Deep
Survey with the Very Large Survey Telescope (Peletier et al., 2020), which is
imaging the Fornax cluster in $ugri$ at comparable depth to the NGVS in Virgo.
Together these surveys will provide the means to identify BCs both through
their young blue stellar populations and, where it exists, their H i gas.

##  9 Conclusions

We have presented follow up HST, VLT/MUSE, VLA (and GBT) observations of five
candidate young, blue, faint, stellar systems in the direction of the Virgo
cluster that are analogous to SECCO 1. With the exception of one spurious
object, we find that these are all comparable to SECCO 1 in terms of their
extremely low stellar masses, blue stellar populations, and high
metallicities, leading us to conclude that they must have formed from gas
stripped from more massive galaxies. However, only one is detected in H i,
suggesting that the others have likely survived sufficiently long to lose much
of their initial gas content. Some of these objects are also surprisingly
isolated, residing several hundred kpc from the nearest potential source of
gas, which poses a challenge for robustly identifying their points of origin.

We considered both tidal and ram pressure stripping scenarios as the potential
formation mechanism of these stellar systems. Although we cannot confidently
exclude either of these mechanisms, and indeed there may not be one single
mechanism responsible for all BCs, ram pressure stripping is most consistent
with the observed properties. In particular, the isolation of some BCs is
difficult to explain with the low velocities ($\leq$300 km s-1) expected for
ejected TDGs, but can more naturally be explained by ram pressure stripping
proceeding at $>$1000 km s-1. In addition, BCs are likely too low mass to be
long-lived TDGs. However, gas clumps formed in ram pressure stripping
simulations are typically much lower mass than BCs (based on the H i masses of
BC3 and SECCO 1), and we suggest that these objects may be atypical and form
from the first loosely bound gas to be stripped, or as a result of stripping
in a clumpy ICM. These massive clumps ($\sim$107 M⊙) of stripped gas moving at
high speed can likely survive sufficiently long in the ICM to form the stellar
populations observed and to become relatively isolated. However, they will
ultimately lose their gas content (the majority of their total mass) and
likely become unbound.

BCs therefore represent a new class of stellar system that form from large
($\sim$107 M⊙) clumps of pre-enriched, stripped gas, are (assumed to be) dark
matter free, and are capable of surviving sufficiently long in the hostile ICM
to become isolated ($>$100 kpc away) from their parent galaxies.

A further census of this class of object in the Virgo cluster, and potentially
the Fornax cluster, will allow for improved constraints on their lifetimes and
how frequently they are produced. However, robust identification of their
parent objects will remain challenging, owing to their isolation. Deep, wide-
field H$\alpha$ imaging, to identify diffuse emission, is a potential approach
for systems where the majority of the neutral gas has already been evaporated.

The authors thank Kyle Artkop for assistance in identifying blue candidates in
Virgo. We also thank Toby Brown and co-authors for providing their X-ray
mosaic from archival ROSAT observations of the Virgo cluster. This work is
based on observations made with the NASA/ESA Hubble Space Telescope, obtained
at the Space Telescope Science Institute, which is operated by the Association
of Universities for Research in Astronomy, Inc., under NASA contract
NAS5-26555. These observations are associated with program # HST-GO-15183.
Support for program # HST-GO-15183 was provided by NASA through a grant from
the Space Telescope Science Institute, which is operated by the Association of
Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.
It is also based on observations collected at the European Organisation for
Astronomical Research in the Southern Hemisphere under ESO programme
0101.B-0376A. This work used both previously unpublished and archival data
from the Karl G. Jansky Very Large Array. The National Radio Astronomy
Observatory is a facility of the National Science Foundation operated under
cooperative agreement by Associated Universities, Inc. The data were observed
as part of programs 13A-028 (PI: J. Cannon) and 18A-185 (PI: K. Spekkens). The
work used images from the Dark Energy Camera Legacy Survey (DECaLS; Proposal
ID 2014B-0404; PIs: David Schlegel and Arjun Dey). Full acknowledgment at
https://www.legacysurvey.org/acknowledgment/. DJS acknowledges support from
NSF grants AST-1821967 and 1813708. MB acknowledges the financial support to
this research from the INAF Main Stream Grant 1.05.01.86.28 assigned to the
program The Smallest Scale of the Hierarchy (SSH) KS acknowledges support from
the Natural Sciences and Engineering Research Council of Canada (NSERC). BMP
is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship
under award AST-2001663. EAKA is supported by the WISE research programme,
which is financed by the Dutch Research Council (NWO). GB acknowledges
financial support through the grant (AEI/FEDER, UE) AYA2017-89076-P, as well
as by the Ministerio de Ciencia, Innovación y Universidades, through the State
Budget and by the Consejería de Economía, Industria, Comercio y Conocimiento
of the Canary Islands Autonomous Community, through the Regional Budget. JS
acknowledges support from the Packard Foundation. MPH acknowledges support
from NSF/AST-1714828 and grants from the Brinson Foundation. JMC, JF, and JLI
are supported by NST/AST 2009894. R. R. M. gratefully acknowledges support by
the ANID BASAL project FB210003. Research by DC is supported by NSF grant
AST-1814208. AK acknowledges financial support from the State Agency for
Research of the Spanish Ministry of Science, Innovation and Universities
through the ”Center of Excellence Severo Ochoa” awarded to the Instituto de
Astrofísica de Andalucía (SEV-2017-0709) and through the grant
POSTDOC_21_00845 financed from the budgetary program 54a Scientific Research
and Innovation of the Economic Transformation, Industry, Knowledge and
Universities Council of the Regional Government of Andalusia.

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##  Appendix A VLA spectrum of BC2

Although the HST imaging of BC2 (Figure 3) indicates that it is a background
group of galaxies rather than a blue stellar system at the distance of Virgo,
we have included an H i spectrum extracted from the VLA data cube for
completeness (Figure 14). As with the other BCs the spectrum is extracted with
a beam-sized aperture centered on the optical position of BC2. However, unlike
the other VLA spectra this spectrum covers a broader range of velocities as
BC2 has no known redshift from H$\alpha$ emission. As expected, no significant
H i signal could be identified.

Figure 14: H i spectrum in the directions of BC2 extracted from the VLA H i
data cube within an aperture equal in area to the synthesized beam.

##  Appendix B Search for points of origin

In the following subsections we search for neighboring galaxies that could
have supplied the gas from which the BCs formed. In particular, we consider
all ALFALFA (Haynes et al., 2018), Virgo Cluster Catalog (VCC, Binggeli et
al., 1985), and Extended Virgo Cluster Catalog (EVCC, Kim et al., 2014)
galaxies in the vicinity of each BC, in an attempt to identify the most
probable point of origin in each case.

###  B.1 BC1

As shown in Figure 12 (top left), the nearest galaxy (in projection) with a
redshift within 500 km s-1 of BC1 ($cz_{\odot}=1118$ km s-1), is AGC 224219
(LSBVCC 79), 14.5′ to the SW, which was detected in H i by ALFALFA and AGES
(Arecibo Galaxy Environment Survey, Taylor et al., 2013, AGESVC2_30), and
marginally detected near the primary beam edge of our VLA observation. After
this there is the large spiral NGC 4579 (VCC 1727) further to the SW (30′ in
total), which was also detected in H i by ALFALFA.

AGC 224219 ($cz_{\odot}=1019$ km s-1) is a very LSB dwarf with a central
surface brightness of 24.8 mag arcsec-2 in $g$-band (Davies et al., 2016). Its
stellar population is also clearly redder than BC1, which comes as a surprise
given that it has at least an order of magnitude more H i. The other
progenitor possibility, NGC 4579 ($cz_{\odot}=1517$ km s-1) is a much larger,
gas-bearing, spiral galaxy. This galaxy was imaged by the VIVA (VLA Imaging of
Virgo spirals in Atomic gas, Chung et al., 2009) survey and they concluded
that it is slightly H i-deficient and its H i distribution is somewhat
truncated (presumably by ram pressure stripping). However, other than this,
the H i distribution and kinematics are regular and there is no sign of an
ongoing or recent interaction. ALMA (Atacama Large Millimeter Array)
observations of its CO molecular gas from VERTICO (Virgo Environment Traced in
CO survey; Brown et al., 2021) present a similarly regular morphology. We also
note that if BC1 originated from NGC 4579 then it must have had a very large
ejection velocity, at least 400 km s-1 along the line-of-sight (LoS). If we
also assume BC1 is on the order of 200 Myr in age (c.f. BC3, Jones et al.,
2022) then the component perpendicular to the LoS would need to be on the
order of 500 km s-1, giving a total ejection velocity of $\sim$650 km s-1, a
large, but not impossible value.

Perhaps both BC1 and AGC 224219 originated from a past stripping event of NGC
4579. They are both along the same direction from NGC 4579 and both have very
similar radial velocities. This scenario would also be consistent with the
observation that AGC 224219 has more H i gas than BC1, as the latter would
have traveled twice as far from its parent object ($\sim$150 Mpc, projected,
in total). However, this scenario would not naturally explain why AGC 224219
appears to have an older, redder stellar population, while BC1 is dominated by
young, blue stars (Figure 2) and H ii regions. Furthermore, if a large
quantity of gas had been stripped from NGC 4579 in the recent past then its H
i and CO distributions would probably show more signs of disturbance and other
clumps of in situ SF would be expected in the immediate vicinity.

The discussion above exhausts the possible sources of gas, within 1∘
($\sim$300 kpc) and $\pm$500 km s-1, from which BC1 could have formed.
Assuming that we have not erroneously rejected the true source of the gas that
formed BC1, we are left with three possibilities: a) the origin of the gas
lies farther away than 1∘, b) the gas originally belonged to a faint or LSB
object that we have missed or excluded due to a lack of a redshift, or c) that
BC1 was a (long-lived) dark object that contained gas but no stars prior to
its recent SF episode.

Based on the simulations of Kapferer et al. (2009), neutral gas being ram
pressure stripped can be found several hundred kpc from the source galaxy.
However, increasing the search radius gives a large number of potential
candidates owing to the density of the Virgo region. If we narrow these down
to those that are blue or contain H i, and are very close in velocity to BC1
then VCC 1686 (IC 3583, $cz_{\odot}=1122$ km s-1, 71′ away), VCC 1992 (IC
3710, $cz_{\odot}=1012$ km s-1, 75′ away), AGC 225847 ($cz_{\odot}=982$ km
s-1, 81′ away), and VCC 1931 ($cz_{\odot}=1183$ km s-1, 83′ away) are the most
promising candidates (note that these are beyond the FoV shown in Figure 12).
These are all irregular dwarf galaxies that appear to be forming stars and are
likely interacting with other galaxies in the cluster.

VCC 1686 can be discounted as a foreground galaxy based on a TRGB distance
measurement (Karachentsev et al., 2014). Estimating the stellar masses of the
remaining 3 candidates (assuming $D=16.5$ Mpc) with the mass-to-light ratio of
Taylor et al. (2011) and photometry of Kim et al. (2014), then applying the
MZR of Andrews & Martini (2013), reveals that all have expected metallicities
at least 0.2 dex lower than BC1 (Table 3). Thus we are left with no strong
candidate that matches the properties of BC1.

###  B.2 BC4 & BC5

There are no galaxies with significant H i reservoirs (detectable in ALFALFA
at the distance of Virgo) within 30′ and 500 km s-1 of BC4 (Figure 12, bottom
left). However, the only (possible) Virgo member in the VCC or EVCC within 30′
is VCC 824 ($cz_{\odot}=392$ km s-1), a red ($g-i=0.84$), nucleated, LSB
dwarf. The nearest galaxy detected in H i is NGC 4419 (VCC 958,
$cz_{\odot}=-275$ km s-1), approximately 40′ to the north of BC4. The H i
distribution of this galaxy has been greatly truncated (Chung et al., 2009),
indicating that it has undergone significant ram pressure stripping. The
VERTICO CO map of NGC 4419 (Brown et al., 2021) indicates that the molecular
gas may extend roughly to the south. This could indicate that this galaxy has
already fallen past the cluster center and is now moving to the north, in
which case this could be a good candidate origin of the gas that formed BC4.
The Taylor et al. (2011) relation for stellar mass (based on $M_{i}$ and
$g-i$) gives an estimate of $\log M_{\ast}/\mathrm{M_{\odot}}=10.25\pm 0.10$
for NGC 4419 (assuming $D=16.5$ Mpc). Using this stellar mass estimate in the
MZR of Andrews & Martini (2013) gives a metallicity estimate of
$12+\log(\mathrm{O/H})=8.74\pm 0.18$, which matches remarkably well with the
observed metallicity of the H ii regions in BC4, $\langle
12+\log(\mathrm{O/H})\rangle=8.73\pm 0.15$ (§4.2).

A little farther to the north there is another gas-bearing galaxy, NGC 4396
(VCC 865, $cz_{\odot}=-122$ km s-1, 77′ away). This galaxy was also studied in
the VIVA project, which found that it is in the process of being stripped by
ram pressure. However, in this case the gas clearly extends to the NW,
indicating that NGC 4396 is unlikely to be the source of the gas that formed
BC4.

Skipping over a few low-mass H i detections (to which we shall return), the
next nearest gas-bearing galaxies are to the south, NGC 4402 (VCC 873) and NGC
4438 (VCC 1043), both of which are $\sim$1.3∘ from BC4 (beyond the FoV in
Figure 12). NGC 4402 was imaged in H i by the VIVA survey, revealing that the
galaxy is deficient in H i, exhibits a truncated disk and clear signs of
ongoing ram pressure stripping, with the current gas tail emanating in a NW
direction (Crowl et al., 2005). However, NGC 4402 has $cz_{\odot}=230$ km s-1,
which places it $\sim$300 km s-1 away from BC4. Although this velocity
separation does not rule out NGC 4402 as the origin of BC4, it means it is
less favored than other candidates. NGC 4438 is just to the east of NGC 4402,
and unfortunately does not have any existing H i imaging. However, the DECaLS
images of NGC 4438 display a highly irregular morphology and the galaxy
appears partially disrupted. The central velocity of the H i emission (Haynes
et al., 2018) is $cz_{\odot}=104$ km s-1, placing it significantly closer to
BC4 in velocity. Again using the Taylor et al. (2011) stellar mass-to-light
relation and the photometry from the EVCC, these two galaxies have stellar
mass estimates of $\log M_{\ast}/\mathrm{M_{\odot}}=9.98\pm 0.10$ and
$10.57\pm 0.10$, respectively. Thus neither would match with the metallicity
of BC4 (via the MZR) quite as well as NGC 4419, but would agree comfortably
within the relation’s scatter.

We also note that BC4 is in the vicinity of SECCO 1 and Sand et al. (2017)
suggested the same sub-group of Virgo (containing NGC 4402 and NGC 4438) as
the potential source of the gas in that object. The extension of NGC 4402 to
the NW (Crowl et al., 2005; Chung et al., 2009) and the highly disturbed
H$\alpha$ emission (Kenney et al., 2008; Sand et al., 2017) in the vicinity of
both galaxies makes them both plausible candidates for the origin of SECCO 1.
However, this does not preclude them also being candidates for BC4. It is
possible that several such objects may be formed from the ram pressure
stripping of a single large galaxy, but the significantly lower metallicity of
SECCO 1 (though still high for such a low mass object,
$12+\log(\mathrm{O/H})=8.38$, Beccari et al., 2017a) suggests that it was
likely formed from a different parent object than BC4.

Finally, we return to the two blue, irregular, dwarf galaxies to the south of
BC4 that we skipped over above, VCC 1001 ($cz_{\odot}=340$ km s-1, 43′ away)
and VCC 945 (IC 3355, $cz_{\odot}=-16$ km s-1, 73′ away). These are likely too
small and far away from BC4 to have been the source of the gas it must have
had recently to support its SF. However, they may be related to BC4, or indeed
to NGC 4402 and NGC 4438. It is also worth noting that three very similar
objects were seen in the vicinity of BC1 on the other side of the cluster.

BC5 is approximately 45′ north of BC4, raising the possibility that perhaps it
is part of the same extended structure, indeed the two objects have almost
identical metallicity measurements (Table 3) and radial velocities (Table 1).
The closest galaxy detected in H i (and within 500 km s-1) is NGC 4419 (VCC
958, $\sim$10′ to the SE), as it was for BC4 (Figure 12, bottom right). Again
this object would be an excellent match for BC5’s metallicity (based on the
MZR). However, if we assume these objects formed from ram pressured stripped
material then it seems unlikely that NGC 4419 could be the origin of both BC4
and BC5, as BC4 is 40′ to its south and BC5 is 10′ to its NW, while ram
pressure tails usually extend in one general direction (i.e. in the wake of
the galaxy’s motion through the ICM). If instead they were formed from tidally
stripped gas then this objection is removed, as they may have formed from
tidal tails on opposite sides of the galaxy.

The next nearest galaxy is NGC 4396, 30′to the NW. As discussed above, this
galaxy is also being stripped, however, as with BC4, the H i tail extends in
the wrong direction.

These two galaxies are the only H i-bearing galaxies within 1∘ and $\pm 500$
km s-1. The only other galaxy in the VCC or EVCC within this range is VCC 583,
a red, LSB dwarf elliptical approximately 58′ to the west. This dwarf is next
to the H i-bearing galaxy NGC 4312, which might have been a good candidate for
the origin of both BC4 and BC5 via ram pressure stripping, except that
multiple TFR distance estimates place it well in the foreground of the cluster
(Yasuda et al., 1997; Russell, 2002; Sorce et al., 2014).

Thus if we wish to look for better-matched candidates then we must look
further afield. Here we come across a promising candidate, UGC 7695 (VCC 1450,
IC 3476; $cz_{\odot}=-159$ km s-1, $\sim$1.5∘ to the SE of BC4, beyond the FoV
shown in Figure 12). This galaxy was extensively studied by Boselli et al.
(2021b). It displays a striking bow-shaped region bright in H$\alpha$ and UV,
a clear sign of ram pressure stripping. Boselli et al. (2021b) argue that this
galaxy is likely being stripped almost edge-on and that the stripped gas is
forced up, over and around the disk. The metallicity of UGC 7695 is also
consistent with that of BC4 and BC5, with the highest values measured by
Boselli et al. (2021b) being $12+\log(\mathrm{O/H})\approx 8.7$ and the
average being $\left<12+\log(\mathrm{O/H})\right>=8.60\pm 0.12$ (Hughes et
al., 2013).

The orientation of UGC 7695 and its velocity relative to the cluster center
suggest that the transverse velocity through the cluster is likely on the
order of 1500 km s-1 (Boselli et al., 2021b). Based on this velocity, the
material that formed BC4 and BC5 would have needed to have been stripped
approximately 250-300 Myr ago. This is consistent with the approximate ages of
the stellar populations of the BCs that we estimated in §4.4, in that those
age estimates are less ($\sim$200 Myr). But, Boselli et al. (2021b) conclude
that the structures seen in UGC 7695 are considerably younger, 50-150 Myr.