# The CO emission in the Taffy Galaxies (UGC 12914/5) at 60pc resolution-I:
The battle for star formation in the turbulent Taffy Bridge

P. N. Appleton  Caltech/IPAC, MC 314-6, 1200 E. California Blvd., Pasadena, CA
91125, USA<EMAIL_ADDRESS>B. Emonts  National Radio Astronomy
Observatory, 520 Edgemont Road, Charlottesville, VA 22903, USA  U. Lisenfeld
Departamento de Fisica Teorica y del Cosmos, Universidad de Granada, Spain and
Instituto Carlos I de Fisica Teorica y Computacional, Universidad de Granada,
Spain  E. Falgarone  LERMA/LRA, Observatoire de Paris, PSL Research
University, CNRS, Sorbonne Universites, UPMC Universite Paris 06, Ecole
normale superieure, 75005 Paris, France  P. Guillard  Sorbonne Universités,
UPMC Paris 6 et CNRS, UMR 7095, Institut d’Astrophysique de Paris, 98 bis Bd
Arago, 75014 Paris, France  F. Boulanger  UPMC Universite Paris 06, Ecole
normale superieure, 75005 Paris, France  J. Braine  Laboratoire
d’Astrophysique de Bordeaux, Univ. Bordeaux, CNRS, B18N, allée Geoffroy Saint-
Hilaire, 33615 Pessac, France  P. Ogle  Space Telescope Science Institute,
3700, San Martin Drive, Baltimore, MD21218, USA  C. Struck  Dept. of Physics
and Astronomy, Iowa State University, Ames, Iowa, 50011  B. Vollmer
Université de Strasbourg, F-67000 Strasbourg, France  T. Yeager  Lawrence
Livermore Laboratories, 7000 East Avenue, Livermore, CA 94550

(Accepted March 28 2022)

###### Abstract

We present ALMA observations at a spatial resolution of 0.2 arcsec (60 pc) of
CO emission from the Taffy galaxies (UGC 12914/5). The observations are
compared with narrow-band Pa$\alpha$, mid-IR, radio continuum and X-ray
imaging, plus optical spectroscopy. The galaxies have undergone a recent head-
on collision, creating a massive gaseous bridge which is known to be highly
turbulent. The bridge contains a complex web of narrow molecular filaments and
clumps. The majority of the filaments are devoid of star formation, and fall
significantly below the Kennicutt-Schmidt relationship for normal galaxies,
especially for the numerous regions undetected in Pa$\alpha$ emission. Within
the loosely connected filaments and clumps of gas we find regions of high
velocity dispersion which appear gravitationally unbound for a wide range of
likely values of $X_{\rm CO}$. Like the “Firecracker” region in the Antennae
system, they would require extremely high external dynamical or thermal
pressure to stop them dissipating rapidly on short crossing timescales of 2-5
Myrs. We suggest that the clouds may be transient structures within a highly
turbulent multi-phase medium which is strongly suppressing star formation.
Despite the overall turbulence in the system, stars seem to have formed in
compact hotspots within a kpc-sized extragalactic HII region, where the
molecular gas has a lower velocity dispersion than elsewhere, and shows
evidence for a collision with an ionized gas cloud. Like the shocked gas in
the Stephan’s Quintet group, the conditions in the Taffy bridge shows how
difficult it is to form stars within a turbulent, multi-phase, gas.

galaxies: interactions — galaxies: individual (UGC 12914, UGC 12915) —
(galaxies:) intergalactic medium

††journal: ApJ

##  1 Introduction

Major mergers between massive gas rich galaxies are transformative events in
galaxy evolution. In many mergers, models suggest that tidal torques within
host galaxy disks drives gas inwards, often forming an intense dust-enshrouded
nuclear starburst (e. g. Mihos & Hernquist 1996). Such a mechanism has been
long accepted as the main explanation for the existence of rare, but powerful
sources of far-IR emission (LIRGs and ULIRGs)111(U)LIRG=(Ultra)Luminous
Infrared Galaxies in the local universe are defined as having high far-IR
luminosities LIR $>$ (1012 L⊙) 1011L⊙. in the local Universe (e.g. Soifer et
al. 1984a, b; Sanders et al. 1986; Armus et al. 1987; Sanders & Mirabel 1996;
Armus et al. 2009; Saito et al. 2015; Armus et al. 2020).

While much of the early modeling of colliding and massive merging galaxies
primarily was aimed at more general collision geometries, head-on collisions,
thought to be responsible for collisional ring galaxies, have always been a
special case (Lynds & Toomre, 1976; Toomre, 1978; Appleton & Struck-Marcell,
1987; Struck-Marcell, 1990; Gerber et al., 1996). Although several of these
systems involved dissipative gas-rich collisons (e. g. Appleton et al. 1996;
Charmandaris & Appleton 1996; Higdon 1996; Braine et al. 2003, 2004), there
has been a resurgence of interest in the treatment of the dissipative effects
of the gas (Renaud et al., 2018), especially the formation of a “splash”
bridge (Struck, 1997; Yeager & Struck, 2019, 2020a, 2020b). Splash gas bridges
are produced when two gas-rich disk systems collide nearly head-on. In such
cases, the stellar components pass through each other, but leave behind a
massive gas bridge (Vollmer et al., 2012). These kinds of collisions are
challenging to models because a large fractions of the gas is strongly
compressed and heated during the collision, and much of the multi-phase medium
remains at various stages of cooling tens of millions of years after the
collision (e. g. Yeager & Struck 2020a, b). These latter models suggest that
the gas between the galaxies continues to collide with high Mach-number,
generating strong turbulent conditions in the bridge. Gas in such bridge
systems can be used to study the effects of turbulence and shocks on the
formation of stars in a relatively “clean” environment, far from the
complicating effects of nuclear starbursts or AGN found in many other major
merger systems.

We present new high-resolution CO observations of one of the best studied
splash bridge systems, known as the Taffy galaxies (UGC 12914/5, Condon et al.
1993; stellar masses of 7.4 and 4.2 $\times\leavevmode\nobreak\
10^{10}\leavevmode\nobreak\ M_{\odot}$ respectively; Appleton et al. 2015).
The two gas rich disks are believed to have collided almost face-on at high
velocity (600-800 km s-1) 25-30 Myr ago (see Figure 1a). The ionized gas in
the galaxy disks appear to be counter-rotating (Joshi et al., 2019),
suggesting the discs has oppositely oriented spins when the disks first
collided. This geometry would create an even stronger affect on the gas
collisions at the time of impact. We are now likely viewing the system almost
edge-on after the stellar components have passed through each other (Vollmer
et al., 2012, 2021), leaving behind a massive neutral and molecular gas bridge
in the center of mass frame of the galaxies.

In addition to the disturbed CO distribution studied with varying degrees of
spatial resolution (Gao et al., 2003; Zhu et al., 2007; Braine et al., 2003;
Vollmer et al., 2021), there is independent evidence that the gas in the
bridge is highly disturbed. S͡pitzer IRS observations have shown the existence
of large amounts of warm (T $\sim$ 100-300 K) emitting molecular hydrogen
(Peterson et al., 2012) with properties consistent with shock or turbulent
heating (Guillard et al., 2009). The spectra showed powerful dominant emission
lines of pure-rotational H2, with large ratios of warm H2/PAH and H2/FIR,
similar to those seen in the Stephan’s Quintet intergalactic shock (Appleton
et al., 2006; Cluver et al., 2010). Herschel PACS observations showed that the
bridge also emits strong [CII]158$\mu$m and [OI]63$\mu$m emission, as well as
emission from [CI], CO (4-3) and CO (5-4) based SPIRE FTS observations
(Peterson et al., 2018). The strength and unusual line ratios of the fine
structure lines point towards heating by shocks and turbulence, as suggested
by recent models (Yeager & Struck, 2020b). Direct evidence for fast atomic
shocks (V $\sim$ 200 km s-1) was found in the bridge (Joshi et al., 2019), as
well as potentially faster shocks from Chandra observations of soft X-ray
emission (Appleton et al., 2015). Radio continuum emission from the bridge
(Condon et al., 1993), can also be explained as a result of Fermi
accelerations of cosmic rays in shocks generated within the turbulent gas
(Lisenfeld & Völk, 2010).

Figure 1: The Taffy galaxy system, a) shown in a false color SDSS bri-band
(blue = b, green = r, red = i) image designed to show the faintest emission,
including a faint extragalactic HII region just south-west of UGC 12915, b)
red circles show the primary beams of the 37 ALMA 12-m pointings (FWHM = 25.6
arcsec) for CO (2-1) observation described here. 1 arcsec corresponds to 300
pc at D = 62 Mpc

Recently Vollmer et al. (2021) has reported the highest spatial resolution
(2.7 arcsec) observations of CO (1-0) in the Taffy system with the IRAM
Plateau de Bure Interferometer (PdBI), as well as detailed models of the
structure and kinematics of the large scale gas distribution. The results
provide very strong support for the collisional picture, and evidence for star
formation suppression in the bridge. The results suggest that much of the gas
in the bridge is not virialized on the 800 pc to few kpc scales, and they
suggest turbulent adiabatic compression is responsible for the high velocity
dispersion in the observed gas clumps. One exception is a luminous
extragalactic HII region (hereafter X-HII region) which may have formed in the
bridge, close to the northern-most Taffy galaxy UGC 12915.

This paper presents Atacama Large Millimeter Array (ALMA)222ALMA, an
international astronomy facility, is a partnership of ESO, the U.S. National
Science Foundation (NSF) and the National Institutes of Natural Sciences
(NINS) of Japan in cooperation with the Republic of Chile. observations of the
CO (2-1) emission from the Taffy system on angular scales (0.24 x 0.18
arcsec2), an order of magnitude higher than previous work.

Because of the large quantity of data obtained in the ALMA mosaics, we split
our discussion of the ALMA CO data into two papers. The current one
concentrates on the CO (2-1) emission from the gas bridge from Cycle 4. A
second paper will discuss more fully the condition of the molecular gas in the
two Taffy galaxies themselves, and will present CO (3-2) data obtained both in
Cycle 4, and in Cycle 7.

The main goals of this paper are, 1) to explore the distribution of molecular
gas in the Taffy bridge at 60-100 pc resolution and its relationship to
Pa$\alpha$ emission observed with HST at similar resolution, 2) probe the
kinematics of the CO gas emission and how that relates to shocks and star
formation previously observed from ground and space based data, and 3) study
in greater detail the one major area of star formation in the bridge (the
extragalactic HII region). This region may provide further insights into the
formation of stars in turbulent environments through comparison with new
optical spectroscopy, new radio continuum, and archival Spitzer and X-ray
observations.

The paper is organized as follows. The observations and data calibrations are
described in §2. Results, including the large-scale molecular distribution and
kinematics and gas surface densities estimates are presented in §3. §4
describes the ionized gas emission in the bridge and in the X-HII regions. §5
is concerned with the relationship between molecular gas and star formation in
the bridge, including testing the Kennicutt-Schmidt relationship for the
clouds, the virial properties of the clouds, and quantifying the star
formation rate in the X-HII region. The Taffy bridge is compared with other
similar intergalactic environments in §6. The possible origin of large-scale
star formation suppression in the bridge is described in §7. The conclusions
are given in §8. An Appendix includes tabulated data and figures relating to
the extracted regions discussed in the main body of the paper.

We assume a distance to the Taffy galaxies of 62 Mpc based on a mean
heliocentric velocity for the system of 4350 km s-1, and a Hubble constant of
70 km s-1 Mpc-1. At this distance, 1 arcsec corresponds to 300 pc.

##  2 Observations

###  2.1 CO Observations and reduction

Table 1 is a complete list of the observations made with the ALMA 12-m arrays
in programs 2016.1.01037.S (40 antennas) and 2019.1.01050.S (41 antennas). In
2016.1.01037.S, 37 full-sampled primary beam pointings with the 12-m array
were made of the Taffy pair and bridge in 12CO$(2-1)$ ($\nu_{rest}$=230.54
GHz; see Figure 1b). Two sets of seven pointings were also made in the
12CO$(3-2)$ ($\nu_{rest}$=345.795 GHz) centered on the brightest part of the
bridge, and a bright region in UGC 12914. A further set of CO (3-2)
observations were made of the bridge and UGC 12915 in program 2019.1.01050.S
(41 antennas). Although we list all the observations in the table, we postpone
discussion of the CO(3-2) observations until a second paper.

The CO (2-1) observations centered on a heliocentric (optically defined)
velocity of 4450 km s-1 were observed with a total bandwidth of 1.875 GHz
(2472 km s-1) and a channel separation of 1.953 MHz (2.6 km s-1) in ALMA Band
6. A second continuum baseband was centered on $\nu$=228 GHz ($\lambda$1.3mm,
and bandwidth 1.875GHz).

Calibration of the CO(2-1) data was performed using ALMA flux and phase-
reference calibration sources during the course of the observations. These
data were processed with a standard ALMA calibration pipeline included in CASA
v.5.5.1-5 resulting in fully flux, phase and bandpass-calibrated visibility
data. The quality of the calibration was carefully reviewed before performing
exploratory Fourier transforms of the ALMA visibility data to produce initial
“dirty channel maps” smoothed to a resolution of 10 km s-1. The maps were made
in each channel over a scale of 8000 x 8000 pixel2, where the pixel scale was
0.018 arcsec.

Table 1: Detail of ALMA Observations  Cycle | Band | Array | Date | Line | Frequency | Mosaic | Total | Time | Final Sythesized | Refs.aa1 = this paper, 2 = Paper II-Appleton et al. (in preparation).
---|---|---|---|---|---|---|---|---|---|---
|  |  |  |  | Sky (GHz) | (number) | Time (min) | on-source (min) | Beam size ($"$) |
4 | 6 | 12m C40-6 | 2016-10-04 | CO(2-1) | 227.1 | 37bb37 primary beam positions were observed, fully sampling the Taffy system; see Figure 1b. | 63.3 | 39.6 | 0.24x0.14 | 1
4 | 6 | 12m C40-6 | 2016-10-05 | CO(2-1) | 227.1 | 37bb37 primary beam positions were observed, fully sampling the Taffy system; see Figure 1b. | 64.2 | 39.6 | 0.24x0.14 | 1
4 | 6 | 12m C40-6 | 2016-10-06 | CO(2-1) | 227.1 | 37bb37 primary beam positions were observed, fully sampling the Taffy system; see Figure 1b. | 72.97 | 39.6 | 0.24x0.14 | 1
4 | 6 | 12m C40-3 | 2016-12-03 | CO(2-1) | 227.1 | 37bb37 primary beam positions were observed, fully sampling the Taffy system; see Figure 1b. | 59.6 | 39.6 | 0.24x0.14 | 1
4 | 7 | 12m C40-5 | 2016-10-27 | CO(3-2) | 340.8 | 2-7ccSeven well-sampled primary beam positions were observed in 2 regions, one centered on the Taffy bridge, and a second in the S-E disk of UGC 12914 (to be described more fully in a second paper). | 96.5 | 53.4 | 0.23x0.18 | 2
7 | 7 | 12m C43-5 | 2021-06-30 | CO (3-2) | 340.8 | 14dd41 antennas, 14 pointings covering bridge and inner regions of UGC 12915 | 197.4 | 108.6 | 0.21x0.18 | 2
Figure 2: (a) CO (1-0) emission contours (Vollmer et al., 2021) taken with
PdBI (beamsize 2.7 x 2.7 arcsec2) superimposed on an SDSS i-band image of the
Taffy system (UGC 12915/4) including the bridge emission, (b) A false color
image of the ALMA total velocity integrated intensity CO (2-1) (227GHz) mosaic
at full-resolution (synthesized beam size 0.28x 0.14 arcsec2), built from 37
separate overlapping primary beam pointings. On this large scale view, the
beam size is too small to display. The arrows show the position of the
extragalactic HII region (circle 3 arcsec, marked as X-HII) on both figures,
and the inset shows the position of the peak in 6 GHz nuclear radio emission
(cross; Appleton et al. in prep.) in both galaxies superimposed on a zoomed-in
image of the CO (2-1) emission. 1 arcsec at D = 62 Mpc corresponds to 300pc.

Extended CO emission was suspected from the galaxies and weaker emission from
the bridge, with emission being present not only at the smallest scales
sampled by the ALMA observations (0.22“x 0.18”) but also on slightly larger
scales. This became clear when we initially tried a conventional CLEAN method
(e. g. Högbom 1974) to deconvolve the “dirty maps” using the interferometric
point-source response, or “dirty beam”, for each channel where emission was
detected. Because of the extended emission, this procedure always led to poor
negative large-scale residual flux (bowls) in some of the residual maps.
Instead, we used the multi-scale CLEAN algorithm (hereafter MSCLEAN; Cornwell,
2008) as implemented in the CASA task ‘tclean’. Unlike conventional CLEAN
methods, which assume that the intrinsic brightness distribution of sources is
made up of points sources (corresponding to a set of zero-scale delta-
functions), the multi-scale clean method allows for both point-sources and
larger scales to be present. Conventional CLEAN deconvolution methods
iteratively build up the source distribution out of delta-functions, by
subtracting the dirty beam from the observed dirty maps (the subtraction is
usually done in the uv-plane). MSCLEAN chooses from a set of smoothed dirty
beams (the point source response convolved with the MSCLEAN scale), and
progressively builds a model of the large-scale flux first, followed by flux
on smaller and smaller scales, until it approximates that of a point source
(zero-scale). This deconvolution of emission on different spatial scales
decreases low-level artifacts caused by the PSF, allowing us to better recover
the extended emission in the image. A full description of the method can be
found in Cornwell (2008). Illustrative examples of its application to cases of
nearby galaxies with extended HI emission are provided by Rich et al. (2008).
In our case, we tested, by trial and error, various scales to optimize the
removal of the negative features seen previously in the residual maps. These
final scales used corresponded to scales of zero (delta function), 6, 12 and
24 pixels. The method was then applied to all the channel maps (10 km s-1
separation, covering a heliocentric optical velocity range from -419 to + 388
km s-1 centered at 4350 km s-1 ) containing emission, leading to a large data
cube of the CO(2-1) emission covering a large part of the Taffy field (see
Figure 1b). As presented in Table 1, the maps resulted in a synthesized beam
with an angular resolution of 0.24 x 0.14 arcsec2, which corresponds to a
projected physical scale of $\sim$60 pc for the Taffy system. The rms noise in
each channel map was $\sim$0.7 mJy beam-1.

###  2.2 HST Pa$\alpha$ Observations

Hydrogen recombination lines, like Pa$\alpha$, trace ionized gas associated
with star formation, ionized shocks and other sources of diffuse ionized gas,
and are commonly used to estimate star formation rates (e. g. Kennicutt 1998a;
Calzetti et al. 2007). In §4 we will discuss how the star formation rates
estimated from the Pa$\alpha$ emission must also take into account strong
contamination from shocked gas (Joshi et al., 2019). These authors did not
find evidence for other sources of diffuse emission (e. g. DIG; see Haffner et
al. 2009) because of the low star formation rates in the system. We use
archival Hubble Space Telescope (HST) Near Infrared Camera and Multi-Object
Spectrometer (NICMOS) observations taken in the F187N and F190N filters, which
were centered on UGC 12915 from the archive. These observations involved a
small mosaic covering an area of 45.7 x 45.3 arcsec2 centered on the galaxy.
We subtracted the continuum images to obtain a Pa$\alpha$ image. The NICMOS-
NIC3 image extends over part of the northern bridge, allowing us to compare
the emission-line image with the CO map. Another NICMOS image of UGC 12914 is
also available, but it does not cover any significant part of the bridge and
is not presented here. The NICMOS observations are obtained on a 0.2 arcsec
pixel scale, which slightly under-sample the PSF at this wavelength (FWHM =
0.25 arcsec at 2$\mu$m). This resolution is comparable to that of the CO (2-1)
ALMA observations.

Since the absolute astrometry of the NICMOS images is known to have
significant uncertainty, we aligned the WCS coordinates of the well-defined
nuclear peak of the galaxy in the F190N filter to that of a 6 GHz radio
continuum image obtained with the Karl G. Jansky Very Large Array (VLA)
(Appleton, in prep.) by re-registering the NICMOS image (a total shift of 0.9
arcsec). This offset was confirmed by finding good agreement between several
compact 6 GHz radio sources, and corresponding bright compact Pa$\alpha$ knots
in the western disk of UGC 12915. Similarly, we were also able to confirm the
new coordinates by comparing the position of two radio hotspots embedded in
the X-HII region with the corresponding knots of star formation at the same
position in the Pa$\alpha$ image. From these tests we believe that the
astrometry in the Pa$\alpha$ image is accurate to $\pm$0.1 arcsec (1/2 NICMOS
pixel).

Pa$\alpha$ flux densities were calculated using the conversion from counts s-1
pixel-1 to erg s-1 cm-2 arcsec-2 evaluated using the relation Fline = 1.054
$\times$ PHOTFLAM $\times$ FWHM (in erg s-1 cm-2) where PHOTFLAM is the
photometric calibration parameter obtained from the observation FITS metadata,
and FWHM is full-width half-maximum of the filter333Based on the metadata
associated with the observations, PHOTFLAM = 3.33 $\times$ 10-18 erg s-1 cm-2
Å-1 DN-1, and the FWHM = 147.6 $\AA$. The range of detected emission has a
surface brightness of 1.5 $\times 10^{-16}<\Sigma(Pa\alpha)<4.3\times
10^{-15}$ erg s-1 cm-2 arcsec-2, with a median SNR of 35. Upper limits were
calculated by 2.5 $\times$ the rms ADU/pixel over sample areas in the vicinity
of the filaments. This upper limit, when converted to surface brightness units
is $\Sigma(Pa\alpha)<4.3\times 10^{-17}$ erg s-1 cm-2 arcsec-2.

###  2.3 Radio Continuum Observations

Radio continuum observations provide another means of estimating star
formation through thermal and non-thermal processes associated with active
star formation regions (Condon, 1992; Murphy et al., 2011). Deep observations
were made as part of a radio polarization study of the Taffy (project 19A-378)
in the A-array of the VLA at 1.4 GHz (L-band) and 6 GHz (C-band) during a 12.6
hrs period in 2019 August 10. Although the main aim of the project was to
measure radio polarization, total intensity maps (Stokes I) were made in the
two bands after processing with the task ’tclean’ (using CASA v5.6.1)
resulting in maps with restored synthesized beams of 0.36 x 0.30 arcsecs2
(C-band) and 1.2 x 1.1 arcsec2 (L-band). The rms noise in each map was
2$\mu$Jy beam-1 and 7$\mu$Jy beam-1 in C and L band, respectively. A more
complete discussion of the radio observations will be provided in a future
paper (Appleton, in preparation).

###  2.4 Chandra X-Ray Observations

We make use in this paper of a Chandra X-ray (0.5-8 kev) point-source image
discussed in greater detail in (Appleton et al., 2015), and made available to
us by those authors. The image was obtained in 2013 from a 39.5 ks exposure
onto the back-illuminated S3 Advanced CCD Imaging Spectrometer (Weisskopf et
al., 2000), and was smoothed with a Gaussian of FWHM 1.5 arcsec to emphasize
the compact structure. In (§4.3) we discuss the relationship between the VLA
radio, HST Pa$\alpha$ and Chandra X-ray observations of the extragalactic HII
region . Because the Chandra X-ray observations had poor absolute astrometry,
we used the VLA radio map of the nucleus of UGC 12914 (a point source as
observed by both instruments) to carefully register the Chandra image 0.5-8
kev image, made available to us by those authors, to the radio position. The
re-registration of the Chandra image to the VLA WCS frame resulted in a shift
of 0.4 arcsec. This decisively shows that the brightest Ultra Luminous X-ray
(ULX) source in Taffy system falls within the envelope of the high surface
brightness Pa$\alpha$ and radio continuum emission from the extragalactic HII
regions. A full discussion of the observations can be found in (Appleton et
al., 2015).

###  2.5 Spitzer archival observations

IR emission is emitted from the dust grains heated by photons from young
massive stars in star-forming regions. In $\S$ 4.2 we provide IR properties of
the extragalactic HII region derived from archival Spitzer observations in the
IRAC (Fazio et al., 2004) 3.6, 4.5, 5.8 and 8$\mu$m bands, and in the MIPS
(Rieke et al., 2004) 24$\mu$m band. Images were dearchived from the Infrared
Science Archive (IRSA) held at IPAC, Caltech. For the IRAC and MIPS images,
the photometry reported in § 4.2 for the X-HII region was obtained by
measuring flux densities in a fixed aperture of radius 4.8 arcsec centered on
the clearly resolved emission from the source. The same aperture was used for
the 24$\mu$m data. However, because of the location of the X-HII region close
to the disk of UGC 12915, care was taken to obtain local background estimates
parallel to the extended disk of the galaxy. This effect was less than a few
percent for the IRAC images, but represented a larger source of uncertainty
for the MIPS 24$\mu$m emission. These uncertainties are reflected in the
photometry reported in Table 2.

###  2.6 Palomar 5 meter Spectroscopic Observations

Observations of the Taffy System were made in moderate seeing conditions
($\sim$1 arcsec), and a 1 arcsec-wide slit with the Double Beam Spectrograph
(DBSP; Rahmer et al. 2012) of the Palomar 5 meter telescope on 2021 January 9.
The 600/1000 grating was used on the red arm of the spectrograph. At H$\alpha$
the spectral resolution was 85 km s-1, and the scale along the slit was 0.29
arcsec/pixel. Flux calibration was performed using short observations of the
white dwarf star G191B2B. The total on-source integration time was 3000 s.
These data were reduced using a Python-based pipeline PypeIt (Prochaska et
al., 2020), which performed bias, dark subtraction, flat field correction
(using dome flats), flux and wavelength calibration using internal lamps.

##  3 Results

###  3.1 Large-scale molecular gas distribution

Early CO (1-0) observations by Gao et al. (2003) of the Taffy system not only
detected gas in the galaxies, but also measured large quantities of molecular
gas in the bridge ($1.4\times 10^{10}\leavevmode\nobreak\ M_{\odot}$, assuming
a standard Galactic N(H2)/ICO conversion factor444assumed Galactic conversion
factor $N(H_{2})/I_{\rm CO}=2\times 10^{20}$ cm-2 (K km s-1)-1. ). However,
this mass is likely to be greatly overestimated (Braine et al., 2003; Zhu et
al., 2007) because there is evidence that a much smaller N(H2)/ICO conversion
factor is appropriate in the bridge. Higher resolution observation of the
system with a beam size 2.7x 2.7 arcsecs2 were obtained by Vollmer et al.
(2021) with the PdBI, and the integrated map is presented in Figure 2a
superimposed over an SDSS image. The observations show the bridge is composed
of partially resolved clumps scattered between the galaxies.

In the current paper our ALMA CO (2-1) observations, which have a spatial
resolution fifty times higher than the BIMA observations, and 12 times that of
the PdBI, are shown in Figure 2b. This full resolution moment-0 map was
constructed in the following way: i) each channel map (of width 10 km s-1) was
spatially smoothed to an effective resolution of 0.4 x 0.4 arcsec2, ii) a
valid mask was constructed of all signal within the smoothed map that was $>$
3.5 sigma above the noise per channel, iii) an integrated (moment-0) map was
then made by applying the masks to the full-resolution channels, and summing
the emission spatially in those channel maps (velocity space) where the mask
indicated valid points above the masking threshold. Regions outside the valid
regions were not summed.

The same mask was also applied to the smoothed version of the cube, resulting
in a smoothed moment-0 map, which is presented in Figure 3a. This figure shows
some of the fainter features better than the full resolution image with 8000 x
8000 pixel2 which is not well represented in a small figure. However, the full
resolution map is used for the majority of the analysis. The ALMA observations
show that many of the features, seen in the lower-resolution PdBI map, appear
as narrowly defined filaments and small clumps.

In addition to the bridge regions, which is the main focus of this paper, most
of the detected ALMA emission from UGC 12914 is seen in narrow structures,
including peculiar narrow ripples of emission along its northwestern disk,
with fainter gas filaments breaking off from the disk into narrow strands
which point almost perpendicularly to the main arc of the emission in that
arm. Narrow dense structure is also seen on the south-eastern part of the same
disk. The region around the nucleus of UGC 12914 contains high surface
brightness emission distributed in a series of loops and spiral filaments (see
inset in Figure 2b). On the other hand UGC 12915, which is more edge-on than
its companion, is dominated by a bright south-eastern curved structure, or
possible tightly wound spiral feature, and a high surface brightness inner
edge-on nuclear disk (see upper inset in Figure 2b). The main galaxy disk
extends to the north-east following the inner optical dust lane where it
breaks up into numerous clumps and extended filaments to the far north-east.
Narrow filaments of gas are also seen extending away from many parts of UGC
12915’s disk in strands to the north. Because of the complexity of the system,
we will mainly concentrate on the Taffy bridge in the current paper, and will
return to a more complete description of the CO emission from the galaxy disks
in Paper II.

The structure of the bridge emission is striking. The region of the bridge
closest to the northern Taffy galaxy is composed of a collection of filaments
and bright clumps of emission, some of which give the appearance of a cone-
shaped structure whose apex lies $\sim$ 10 arcsec south-west (3kpc) from the
center of UGC 12915. However, other fainter filaments cross the structure
(Figure 3), and we will show that the velocity structure of the bridge gas is
quite complex, and the cone-shaped pattern does not form a single coherent
kinematic structure. The faint X-HII region, seen in the optical image in
Figure 2a, lies buried in the tangled north-western part of the main emission
from the bridge.

There are many other scattered CO complexes apparent in the observations,
including gas to the far north-west of UGC 12915. Some of this gas may be part
of a separate “north-western” bridge of emission identified in optical
Integral Field Unit (IFU) observations of ionized gas in the Taffy bridge
(Joshi et al., 2019). Previous observations of both the neutral hydrogen
(Condon et al., 1993) and Herschel [CII]158$\mu$m emission (Peterson et al.,
2018) show neutral gas in this area.

An important question to ask is what fraction of the CO emission in the Taffy
bridge is detected on the small (60 - 100 pc) scale compared with single-dish
observations (Braine et al., 2003; Zhu et al., 2007) or large-beam
interferometric observations like BIMA (Gao et al., 2003). According to Zhu et
al. (2007) and Gao et al. (2003), the BIMA observations of the Taffy, made
with a 9.8 x 9.7 arcsec2 beam, capture most of the flux seen in the single
dish observations. Based on a data cube provided by Y. Gao (personal
communication), we centered a 20.5 arcsec diameter circular aperture on the
main concentration in the northern bridge, and estimated the integrated flux
over the measured profile to be 92.1 Jy km s-1 in CO (1-0), which agrees with
a statement made in Gao et al. (2003) of emission associated with the region
they called the “HII region”, which actually includes most of the structures
we have been discussing. To compare this with the ALMA observations requires
converting this flux to an equivalent CO (2-1) flux over the same area.
Assuming r21 = 0.79 (Zhu et al., 2007) (where r21 is the ratio of
ICO(2-1)/ICO(1-0)), the equivalent CO (2-1) flux should be 291 Jy km s-1 for
the same aperture. From our ALMA observations, we derived, by integrating the
CO emission over the same area, a total flux of 161.6 Jy-km s-1. From this we
estimate that with ALMA we detect 55.5$\%$ of the emission seen in the BIMA
observations. Thus a significant fraction of the CO emission in the northern
bridge is in an extended component not detected by ALMA.

###  3.2 Large-scale molecular gas kinematics

In Figure 3b, we show an intensity-weighted (1st Moment) mean velocity map of
the whole system (relative to the average heliocentric velocity of 4350 km
s-1for the two galaxies) with the same 0.4 x 0.4 arcsec2 smoothing as Figure
3a. For reference, the systemic velocities of UGC 12915 and UGC 12914 are
quite similar (-14 and +21 km s-1 respectively, implying that most of the
radial motion between the galaxies is in the plane of the sky (Condon et al.,
1993). In this edge-on view of the collision, the counter-rotation of the two
galaxies is particularly obvious. UGC 12915 appears to show blue-shifted
emission in the SE and red-shifted emission in the NW, whereas UGC 12914 shows
the opposite behavior. This counter-rotation may have contributed to an
increased amount of cloud-cloud collisions when the two galaxies originally
collided, almost face-on (Vollmer et al., 2012; Yeager & Struck, 2019, 2020a,
2020b). The extension of scattered gas clouds to the north-west of UGC 12915,
discussed earlier, shows peculiar motions not consistent with regular rotation
within UGC 12915. While the regular rotation of UGC 12915 is from blueshifted
(-340 km s-1; blue color) to redshifted gas (+320 km s-1; yellow color) in the
figure, the clouds further north and to the west show systemic velocities
between -50 to 100 km s-1 (dark green and green). Similarly, non-corrotating
emission was noted in the velocity field of the ionized gas in this region,
where even more discrepant velocities were observed (Joshi et al., 2019). This
supports the idea that it may be part of a second, kinematically-distinct
bridge between the two galaxies, or the remnants of a tidal tail from UGC
12914 (Vollmer et al., 2021).

In the main CO bridge, the average kinematics is quite complex, as Figure 4a
shows (displaying a velocity range from -100 to 140 km/s). The bridge as a
whole does not show large-scale coherent motion, but is rather made up of many
clumps and filaments with disparate average velocities. Some of the filaments
in the cone-shaped structure do show weak systematic motions along parts of
their length, but, except for the filaments near UGC 12915, appear
kinematically distinct and do not seem obviously related to each other.

Figure 4b shows a representation of the CO (2-1) line width as a function of
position in the bridge. The color coding is based on the FWHM (in km s-1) for
the regions presented in Table 3. The points are superimposed on a contour map
of the total intensity of the CO (2-1) emission. Blue and green points show
the gas with the lowest line-width, whereas many regions have FWHM greater
than 65 km s-1(orange and dark red filled circles), with values extending up
to 115 km s-1. Regions of high velocity dispersion are scattered throughout
the filaments and clumps, with the quiescent gas (FWHM $<$ 40 km s-1) being
the minority. Regions with the lowest velocity dispersion are mainly found in
some of the filaments close to UGC 12915, and in the area near the X-HII
region. Some example spectra are shown in Figure 15.

Although it is often traditional to show channel maps (intensity maps of the
line emission as a function of velocity) of the full velocity cube of the
observations, we will defer this to the second paper, since the emphasis of
the current paper is the star formation properties of the bridge.

Figure 3: (a) Smoothed (to 0.4 x 0.4 arcsec2) image of the CO (2-1) total
velocity-integrated intensity map showing fainter features. The color stretch
encompasses emission as faint as 0.05 Jy km s-1beam-1 (dark blue) to the
brightest at 2 Jy kms s-1beam-1 (white), and (b) intensity weighted mean
heliocentric velocity (optical velocity definition relative to 4350 km s-1) of
the full field.  Figure 4: (a) A zoom-in on the mean velocity field
(heliocentric velocity relative to 4350 km s-1) of the northern bridge to
emphasize more finely the velocity structure. White contours show the
Pa$\alpha$ emission from the X-HII region. (b) A color-coded representation of
the fitted FWHM (km s-1) of the CO-emitting clouds on 85pc scales superimposed
on contours of CO surface density. Green and blue filled circles show clouds
low dispersion, whereas orange and red show higher values. Only measurements
with high quality profiles (Spectral Quality 1.0, see Table 3) are shown.

###  3.3 Molecular surface density in selected regions

We extract the spectra of more than 239 small regions of the CO emission
within the filaments and clumps, and compare their properties to star
formation rate estimates derived from the NICMOS observations. In Figure 5,
the CO bridge filaments and cloud complexes are divided into 12 large regions
which represent areas that seem spatially and kinematically related. Regions A
to G cover emission complexes associated with the brighter part of the
northern bridge. Region D includes emission associated with the X-HII region.
We also including two coherent filaments, H and I, that fall close to the disk
of UGC 12915 and may be part of the bridge. Regions J, to the NW, and the
possible extension of filament A, called AE1, and two bright bridge clumps AE2
are also analyzed. These bridge regions lie outside the area sampled by
NICMOS, but are included for completeness.

Figure 5: The CO (2-1) integrated bridge emission divided into regions of
bright emission. Regions A to G include the filaments and clumps associated
with the main bridge (where D includes emission associated with the X-HII
region) and H and I show filaments that may be part of the bridge but lie
close to UGC 12915. Regions J and AE1/2 are scattered regions in the bridge.
AE1 may be an extension of filament A. Each region is divided further into
many small extracted apertures, the properties of which are shown both
graphically and in tabular form in the Appendix (Figures 16 and 17, and Table
3).

To study the details of the emission in each of the regions shown in Figure 5,
we split each structure up into small extraction aperture sub-regions which
are described more fully in Appendix A, and Figures 16 and 17. Using the
software package CASA and the CASA Viewer, we extracted spectra of each region
(of dimension 0.23 x 0.28 arcsecs2), slightly larger than the resolution of
the ALMA data, and large enough to sample a significant part of the NICMOS
PSF. Each of the extracted spectra were then fit with a Gaussian line profile.
In a few positions (Regions A1 through A4), we observed double-line profiles
along the same line of sight. Here we fit two Gaussian components. For all the
CO spectra, we estimate the mean radial velocity, the FWHM and peak flux
density and finally the integrated CO flux $S_{\rm CO(2-1)}\Delta V$ in Jykm
s-1. The line properties of all the extracted regions are presented in Table
3. Table 3 provides a flag of the quality of each spectrum. Of all the spectra
extracted, 219 were deemed to be of sufficient quality (good baselines, signal
to noise ratio) to be included in the analysis.

To estimate the molecular gas properties from the observed cloud line
properties, it is necessary to make several assumptions, including a
conversion to H2 column density. We will use the standard conversion to
molecular gas mass, including a 36$\%$ correction for Helium (Bolatto et al.,
2013):

$\left(\frac{M_{\rm gas}}{M_{\odot}}\right)=1.05\times 10^{4}X_{\rm
CO,20}S_{\rm CO(1-0)}\Delta VD_{L}^{2}(1+z)^{-1},$ (1)

where $D_{L}^{2}$ is the luminosity distance in Mpc and $X_{\rm CO,20}$ is the
conversion factor $N({\rm H_{2}})/I_{\rm CO}$ in units of 2$\times
10^{20}$cm-2 (K km s${}^{-1})^{-1}$. Here N(H2) is the H2 molecular column
density and ICO is the velocity integrated intensity of the CO (1-0)
transition in K km s-1. Therefore, to derive $M_{\rm gas}$, we also need to
make assumptions about both the value of $X_{\rm CO}$, and the ratio of
$S_{\rm CO(1-0)}/S_{\rm CO(2-1)}$.

Braine et al. (2003) suggested that $X_{\rm CO}$ was probably at least a
factor of 4 times lower in the bridge than the Galactic value, based on single
dish observation of the ratio of the 13CO/12CO (for the 1-0 transition), which
implied the 12CO line was almost optically thin. A similar conclusion was
reached by Zhu et al. 2007, using the transitions CO (3-2), (2-1) and (1-0),
and performing LVG modeling (Goldreich & Kwan 1974). They estimated that in
the bridge $X_{\rm CO}$ was between 2-3.6 $\times 10^{19}$ cm-2 (K km
s${}^{-1})^{-1}$, which is 5 to 10 times lower than $X_{\rm CO,20}$. Both of
these measurements were made with large filled apertures on scales of
$\sim$11-12 arcsec. Recently, Vollmer et al. (2021) adopted an intermediate
value of $X_{\rm CO}$ = 1/3 $X_{\rm CO,20}$ for the bridge in their PdBI
$\sim$ 3 arcsec resolution beam. Given the uncertainty in the LVG modeling,
and the large difference in scale between the previous single dish
observations and our ALMA observations, we adopt as an initial working
hypothesis the Braine et al. (2003) value of $X_{\rm CO}$ = 5$\times 10^{19}$
cm-2 (K km s${}^{-1})^{-1}$, or 1/4 $X_{\rm CO,20}$. We will explore the
implications of varying this value on the derived properties of the clouds and
their line of sight extinction under different assumptions. We finally assume
$S_{\rm CO(1-0)}/S_{\rm CO(2-1)}=(\nu_{1-0}/\nu_{2-1})^{2}$(r${}_{21})^{-1}$,
and r21 = 0.79 (see Zhu et al. 2007). Gas masses from our extracted regions
range from 0.1-1 $\times$ 107 ($X_{\rm CO}/X_{\rm CO,20}$) M⊙.

##  4 Ionized gas emission in the bridge and X-HII region

###  4.1 The Pa$\alpha$ distribution in the bridge

In Figure 6 we show the complex filamentary structure of part of the bridge in
CO (2-1) superimposed on the Pa$\alpha$ NICMOS image of UGC 12915, which
includes the brighter and most interesting parts of the bridge region. The
figure shows that within UGC 12915, the CO is largely confined within the
inner disk, and is well correlated with the bright Pa$\alpha$ emission. A
narrow nuclear CO disk is detected, and there are several CO filaments
extending away from the disk to the north of UGC 12915, which, except for a
few isolated cases, are devoid of obvious star formation.

Figure 6: Contours of the full-resolution ALMA CO (2-1) velocity integrated
intensity superimposed on the false-color image of HST NICMOS image of
Pa$\alpha$ around UGC 12915 including parts of the bridge. Contour levels are
0.1, 0.15, 0.25, 0.3, 0.35, 0.45, 0.55, 0.65, 0.75, 1.0, 1.2, 1.4, 1.75, 1.85
and 1.95 Jykm s-1beam-1.

Concentrating on the bridge region, the most prominent Pa$\alpha$ emission
comes from the bright extended X-HII region (Bushouse, 1987; Jarrett et al.,
1999; Joshi et al., 2019). Several clumps of CO emission are seen projected
against this star forming region, including an elongated finger of faint CO
emission which crosses its center. Very little Pa$\alpha$ emission is seen
from the other CO structures, except for one or two faint possible
associations. The lack of obvious star formation in these dense clumps is a
characteristic of the CO emission observed in the bridge (see also Vollmer et
al. 2021).

For those regions with extracted CO spectra that fall within the area covered
by NICMOS, we then proceeded to extract Pa$\alpha$ flux surface densities of
the emission using the software package SAOImage-DS9555SAOImage DS9
development has been made possible by funding from the Chandra X-ray Science
Center (CXC), the High Energy Astrophysics Science Archive Center (HEASARC)
and the JWST Mission office at Space Telescope Science Institute (Joye &
Mandel, 2003).. Of the 219 high quality CO(2-1) extracted bridge spectra, 24
lay outside the NICMOS field of view either in the southern bridge or to the
NW of UGC 12915. Of the remaining 194 spectra, only 35 (18$\%$) show a
detectable Pa$\alpha$ emission in the bridge, and 5 are associated with Region
H, which is very close to the disk of UGC 12915. Except for the regions B18,
19, 20, 31, 32 and 33, all of the other detected regions are associated
directly with the X-HII region. Elsewhere, only upper limits were obtained for
the surface brightness of the Pa$\alpha$. The flux surface densities and upper
limits are tabulated in Table 3 in units of erg s-1 cm-2 arcsec-2 for the same
extraction regions as those measured for the CO emission line fluxes.

Figure 7: Taffy X-HII region a) Pa$\alpha$ false color image with contours
(blue) of CO (2-1) integrated emission overlaid. The green circles show
photometric apertures (ALL, hotspot N and S) labelled in Table 2. b) 6cm radio
continuum false-color image from VLA A-array (Appleton et al. in preparation)
with contours Pa$\alpha$ emission (red) and the 0.5-8keV X-ray ULX source CXOU
J0001409+232938 (white contours; Appleton et al. 2015), c) same radio images
as in b) but overlaying contours of the Spitzer IRAC band 1 (3.6$\mu$m
emission (white). Figure 8: (a) A position-velocity diagram taken along the
main CO bar-like structure (PA = -8 deg) seen in the integrated map of the
molecular gas. The diagram is reminiscent of two sides of a flat rotation
curve displaced around the position of the brightest 6cm radio source
(Hotspot-S of Figure 7). The gas is extended over a scale of 2 arcsec (600pc),
(b) Line emission taken along the same position angle with the Palomar 5-m
DBSP (slit width = 1arcsec, seeing 1 arcsec) showing the distribution of [NII]
and H$\alpha$ along the slit. The main emission shows a tilted distribution
consistent with a rotating disk of $\pm$ 50-60 km s-1in the inner part of the
structure. The rotation and magnitude is consistent with the putative CO disk.
A blueshifted component is seen in H$\alpha$, but is not detected in CO.

###  4.2 The extragalactic HII (X-HII) region and its ionized gas

Figure 7 shows and compares the X-HII region at different wavelengths. The
overall scale of the Pa$\alpha$ emission from the X-HII region, shown in
Figure 7a, is 2 arcsec (or 600 pc) and is composed of two compact regions of
emission surrounded by more diffuse gas. The CO is dominated by a bar-like
structure with a position angle of -8 degrees (north through east) crossing
the face of the Pa$\alpha$ emission and clumpy structures surrounding it. In
Figure 7b we show the $\lambda$6 cm radio continuum image obtained with the
VLA (Appleton et al., in preparation) which mimics the Pa$\alpha$ structure
(red contours), again showing two dominant emission regions and diffuse
emission. Also shown are (white) contours of X-ray emission defining the
brightest ULX X-ray source CXOU J0001409+232938. The ULX source falls close to
the northern compact region. Figure 7c shows the 6cm radio emission image with
contours of Spitzer IRAC 3.6$\mu$m emission superimposed. The X-HII region is
detected in all four IRAC bands and at 24$\mu$m with MIPS. The centroid of the
3.6$\mu$m image (which has a much lower spatial resolution than the other
images) falls close to the brighter southern compact radio source within the
X-HII region.

Returning to the relationship between the ionized gas and the molecular gas
Figure 8a shows a position-velocity diagram constructed along a position-angle
of the main CO-bar at PA = -8 degrees. This is also the angle projected onto
the sky where the CO emission implies an approximate major axis of rotation.
The resulting position-velocity diagram seems to show evidence of bulk motions
from the north to the south, with an inflection point at the position of
Hotspot-S (Figure 7a), which is best defined by the radio position of the
southern 6 GHz source. Relative to an assumed systemic velocity of 106 km s-1
for the X-HII region666The velocities are all relative to the Taffy mean
heliocentric velocity of 4350 km s-1, the gas shows a sudden jump from
positive velocities (+60 km s-1) to negative ones (-60 km s-1) at the position
of the radio source. This supports the idea that Hotspot-S is the kinematic
center about which the gas is rotating. The scale of this putative disk has a
radius of 1-1.5 arcsec = 300-450 pc at D = 62 Mpc.

To further support the idea that we are observing rotational motion, Figure 8b
shows a part of the Palomar spectrum (slit width = 1 arcsec) positioned along
PA = -8 degrees. Bright H$\alpha$ emission is detected, as well as the
satellite lines of [NII] which are well separated from H$\alpha$. The
H$\alpha$ emission shows two features of note. Firstly the main core of the
emission centroid is tilted, consistent with a rotating bright disk, with
rotation of $\pm$ 100 km s-1across 1-1.5 arcsecs. Even though the Palomar
spectral resolution at H$\alpha$ is much lower (85 km s-1) than the CO data,
its rotational motions appear to have the same sign and approximate magnitude
as the CO emission. This suggests that both the CO and H$\alpha$ are part of
the same disk. We also observe a second, broad line-width blue-shifted
H$\alpha$ component observed at negative velocities, which is centered at the
same position as the main tilted disk-like structure, but shifted by 200-300
km s-1 from it. CO emission is clearly absent in Figure 8a at this velocity.
It is unlikely that the negative-velocity gas represents a powerful outflow
from the X-HII region, since its overall star formation rate, as shown in this
paper (§5.3), is quite low. Discrepant velocity structures, similar to this
one, are seen in the kinematics of molecular gas in the Antennae overlap
region (Tsuge et al., 2021a, b), where there is evidence of much stronger star
formation. They may be evidence for shocks associated with cloud-collisions.
Similar examples of very faint star formation embedded in shocked
intergalactic emission-line gas exhibiting multiple line profiles is observed
in Stephan’s Quintet (see for example Xu+ 2003; Konstantopoulos et al. 2014;
Guillard et al. 2021).

##  5  Relationship between molecular gas and star formation in the bridge

###  5.1 Deriving star formation rates from Pa$\alpha$ observations

In normal galactic disks (e. g. Calzetti et al. 2007, 2010), it is generally
assumed that a measure of the star formation rate can be inferred from the
strength of a relatively unobscured hydrogen recombination line, like
Pa$\alpha$, with the assumption that most of the emission is from gas heated
by Lyman-continuum photons from hot stars.

####  5.1.1 Correcting for shocked gas

In the Taffy bridge, we already have strong evidence from previous studies
(Joshi et al., 2019) that a significant fraction (up to approximately 45$\%$
of H$\alpha$ emission) comes from gas heated in shocks with characteristic
shock velocities of 200-300 km s-1. This result applies not only to the bright
extragalactic HII region but also to fainter extended H$\alpha$ emission seen
throughout the bridge. Ubiquitous shocks throughout the bridge are also
supported by evidence for significant quantities of warm molecular gas in the
bridge which presence cannot be explained by heating by star formation
(Peterson et al., 2012, 2018). Thus, it cannot be assumed that all the
Pa$\alpha$ emission in the Taffy bridge is due to star formation alone. As a
result of this uncertainty, we consider two limiting cases in our study. In
the first, we assume all the Pa$\alpha$ is from star formation. In the second,
we assume a more realistic scenario where only 55$\%$ of the Pa$\alpha$
emission arises from star formation based on the observation of Joshi et al.
(2019). This range of possible values will be reflected in our plots of star
formation surface density in the subsequent discussion.

####  5.1.2 Correcting for extinction

A second source of uncertainty in measuring the star formation surface density
from hydrogen recombination lines is extinction. Given that the calculated H2
surface densities approach 1000 $M_{\odot}$ pc-2 under some assumptions, we
cannot assume that extinction is negligible, even at the wavelength of the
Pa$\alpha$ transition ($\sim$ 1.8$\mu$m). Assuming we can correct for the
shocked fraction, and extinction, we can convert the corrected Pa$\alpha$
surface density fluxes into equivalent H$\alpha$ fluxes, assuming Case B
recombination ($f(H\alpha)/f(Pa\alpha)$ = 8.59; Osterbrock 1989), and then
convert the fluxes to star formation rates via the transformation of Kennicutt
(1998b).

To correct for Pa$\alpha$ extinction both the detected fluxes and upper limits
(which applies to the majority of the clouds) we adopt two different
approaches:

  * •

Case1: A minimal extinction assumption at Pa$\alpha$ based on the visible
light measurements of the Balmer decrement. Low extinction at Pa$\alpha$ can
be inferred from measured values of AV using IFU spectroscopy to estimate the
Balmer decrement across faint ionized gas filaments in the bridge from Joshi
et al. (2019). In the region of the bridge these authors estimated AV $\sim$
1.5 mag. Assuming a Calzetti et al. (1994) extinction law, this would convert
to APaα of $\sim$ 0.2 mag777Adopting other commonly used extinction laws does
not affect this conclusion significant since most extinction curves deviate
from one another much more in the UV, and differ by only a few percent in the
near IR (Gordon et al., 2003). These results are also consistent with our
Palomar 5-m spectroscopy obtained with a 1 arcsec slit size.

  * •

Case 2: Calculate AV from the measured molecular gas surface density: The
estimate of extinction from nebulae emission lines might be biased to lines of
sight with lower extinction. We therefore consider a complementary method and
derive the extinction from the total column density, N(H) (= 2N(H2)) within
each cloud. We assume that H2 dominates the hydrogen column density on the
scale of 0.2 arcsec (e. g. Kahre et al. 2018). This conversion is complicated
by two factors. Firstly, the measurement of the total hydrogen column density
N(H) (= 2N(H2)) depends on our assumed XCO factor. Secondly, the assumed gas-
to-dust mass ratio and the extinction law affect the extinction per total gas
column density, $A_{V}/N(H)$. Neither of these factors are well constrained in
the case of the Taffy CO bridge gas. We will perform a limited exploration of
the effects of changing some of the assumptions.

A knowledge of the gas-to-dust mass ratio, which scales linearly with AV, is
particularly important in the Taffy bridge. Previous observations have
suggested that dust is strongly depleted there. Dust was detected in emission
at 450 and 850$\mu$m using the James Clark Maxwell Telescope (JCMT) with an
angular resolution of 9.4 and 16 arcsec, respectively, by Zhu et al. (2007).
These results suggest unusually large gas-to-dust mass ratios of 600-800 in
the bridge, many factors ($>$5) higher than values for the Galaxy or nearby
galaxies($\sim$ 80-150). These large values are in contrast to the Taffy
galaxies which show more normal gas-to-dust ratios. Dust depletion might be a
result of grain destruction in shocks in the bridge (Jones et al., 1996).
Based on these observational results, we consider a range of dust depletions
from no-depletion, to depletions of up to a factor of 5.

We take the relationship of Kahre et al. (2018) of N(H)/E(B-V) = 5.8 $\times
10^{21}$ H cm-2 mag-1, suitable for Galactic dust, and divide the extinction
by an assumed dust depletion factor to account for the fewer dust particles
per H-atom compared with normal Galactic gas. We also consider other
relationships between AV and and N(H), such as lowering the metallicity of the
gas.

Figure 9: Comparing the KS relation with the molecular and SF properties of
the Taffy bridge under different assumed values for the CO conversion factor
$X_{\rm CO}$ for minimal Pa$\alpha$ extinction (Case 1, see § 5.1.2), a)
assuming $X_{\rm CO}$ = 1/4 $X_{\rm CO,20}$, and b) $X_{\rm CO}$ = 1/2 $X_{\rm
CO,20}$. Grey symbols show the star formation rate density assuming no shock
contribution to Pa$\alpha$, whereas the colored connected symbols are the more
realistic case where shocks contribute up to 45$\%$ of the bridge Pa$\alpha$
emission (Joshi et al., 2019). In each plot, the blue (or grey) downward
arrows show SFR density limits derived from the undetected Pa$\alpha$
emission, with (and without) a shock contribution. The different detected
regions are shown with symbols indicated in the central panel. The blue and
yellow lines shows two representations of the average KS relationship, for
spiral and starburst galaxies (Kennicutt et al., 2007; de los Reyes &
Kennicutt, 2019). The 1$\sigma$ spread is shown as dotted lines. The pale blue
points (and density contours) are individual spatially-resolved CO
measurements made with ALMA across the face of the nearby M51 for reference
(for $X_{\rm CO}$ = ($X_{\rm CO,20})$) (Leroy et al., 2017, 2018).  Figure 10:
The (logarithm of) SFR surface density in the bridge clouds versus molecular
gas surface density, in which the star formation is corrected for extinction
using the hydrogen column density method (Case 2, see § 5.1.2), and for
different assumed values of $X_{\rm CO}$, (a) $X_{\rm CO}$ = 1/4 $X_{\rm
CO,20}$, assuming Milky Way (MW) extinction law, with a dust depletion of a
factor of 5, (b) same as (a) but with $X_{\rm CO}$ = 1/2 $X_{\rm CO,20}$, (c)
same as (a), but with no extra depletion and, (d) $X_{\rm CO}$ = 1/4 $X_{\rm
CO,20}$, assuming an Small Magellanic Cloud(SMC)-like extinction law and no
additional depletion. Significant bulk dust depletion of $\geq$ 5 was measured
in the bridge by Zhu et al. (2007). The symbols and lines in the plot are the
same as for Figure 9.

Table 3 provides surface flux densities and logarithmic star formation rate
surface densities for just one limiting case: that of minimal Pa$\alpha$
extinction and assuming 100$\%$ of the emission arises from star formation. In
the figures and discussion that follows, we will explore a much wider range of
possibilities. The table also provides a Pa$\alpha$ flag that indicates
whether Pa$\alpha$ is an upper limit, a detection, or is outside the NICMOS
field of view.

###  5.2 Testing the Kennicutt-Schmidt Relationship in the Taffy Bridge

Figure 11: The velocity dispersion of the CO-emitting clouds versus the gas
surface density as a diagnostic of cloud stability for two CO/H2 conversion
factors, (a) $X_{\rm CO}$ = 1/4 $X_{\rm CO,20}$ and (b), $X_{\rm CO}$ = 1/2
$X_{\rm CO,20}$. The blue line shows the mean relationship for CO emitting
clouds found by Leroy et al. (2016) from the PHANGS survey of normal galaxies
and the grey shaded zone show the virial stability line for spherical
constant-density clouds with diameters ranging from 85 to 100 pc (see text).
Clouds below the grey shaded zone are gravitationally bound, whereas clouds
above it are unbound.

Figure 9a-b shows a series of logarithmic plots of the star formation rate
surface density versus the gas surface density for two plausible cases of XCO,
and assuming Case 1 extinction (§4.2). In normal galaxies, these two
properties are related in what is referred to as the Kennicutt-Schmidt
(hereafter KS) relationship (Kennicutt, 1998a). Each plot shows the range of
possible star formation density associated with different regions (see key)
for the case where there are no shocks contributing to the Pa$\alpha$ (grey
symbols) connected to points with 45$\%$ shock contamination (colored
symbols). A grey line joins the two extremes.

The figures show that generally the points inside the X-HII region lie close
to the expected KS relationship, but the majority of the other regions do not.
Those with weaker assumed star formation are spread into and below the region
occupied by clouds in M51 from Leroy et al. (2017, 2018). The upper limits
(with or without shock contamination), which include the majority of the more
than 200 CO-emitting regions, lie significantly below the surface density they
would have if they followed the KS relationship. Depending on the assumed XCO
value, the denser clouds lie at least an order of magnitude below the standard
relationship. This strongly suggests these molecular regions show significant
star formation suppression.

Could this conclusion be driven by an underestimation of the actual extinction
in the dense molecular gas? In Case 1 extinction, we assume that the
extinction measured for the X-HII region can be universally applied to the
regions with Pa$\alpha$ upper limits. To explore a more general way of
measuring extinction, independent of optical/near-IR emission line
corrections, we apply Case 2 extinction to all the points, even the upper
limits, in Figure 10 a-d. Figure 10a and b, both with a dust depletion factor
of 5, show the effect of increasing the surface density of the gas from
$X_{\rm CO}$ = 1/4 $X_{\rm CO,20}$ to $X_{\rm CO}$ = 1/2 $X_{\rm CO,20}$,
respectively. Figure 10c shows the (unlikely) case of no dust depletion, and a
Milky Way-type extinction law, whereas Figure 10d shows an SMC-like extinction
law, again with no extra dust depletion (Bohlin et al., 1978; Gordon et al.,
2003)). The symbols and meaning are the same as for Figure 9. As one might
expect, the SMC extinction law looks very similar to the low MW gas-to-dust
ratio case. This latter example is provided to explore the unlikely
possibility that the gas is of low-metallicity–not expected in a head-on
collision of this kind with two massive galaxies.

The results of varying the assumed degree of extinction over a wide variety of
reasonable parameter space (given the measured values for both lower-than-
normal XCO and large dust depletion inferred by Zhu et al. (2007)) continue to
support the idea that the majority of the CO-detected clouds in the Taffy
bridge are strongly star-formation suppressed (Vollmer et al., 2021), even at
the newly explored scale of 60-100 pc relation. The behavior in the diagrams
is similar for all cases. Increasing the column density pushes the cloud
points to both the right (increased H2 surface density) and upwards (more dust
and assumed extinction). The extinction “knob” can be turned by either keeping
the dust depletion constant and increasing the gas surface density, or by
keeping the surface density constant and increasing the dust content, or both.

The plots in Figure 9 and Figure 10 also demonstrate another interesting
result. The sources with detected Pa$\alpha$ associated with the X-HII region
lie generally at, or above the mean KS relationship, depending on the assumed
star formation/shocked gas emission associated with the Pa$\alpha$. A subset
of the explored parameter space would lead the X-HII region exhibits “normal”
star formation rates and high shock contamination with values of $X_{\rm
CO}\sim 1/2X_{\rm CO,20}$ (e. g. Figure 9b; Case 1 extinction, or Figure 10b ;
Case 2 extinction).

In most of the realistic scenarios that fit most of the known properties of
the Taffy bridge, we conclude that there is strong evidence for star formation
suppression in the majority of the molecular gas in the bridge on scales of
60-100 pc. Only in the region where the CO emission is seen projected against
the stronger Pa$\alpha$ emission do we begin to see much more normal
relationship between the surface density of the gas and that of the young
stars.

###  5.3 Small-scale molecular gas kinematics

We next explore the line-of-sight velocity dispersion of the gas as a function
of the gas surface density and its relationship to star formation. The
condition for virial stability is given as $2T/U=1$, where $T$ is the kinetic
energy of the clouds (assumed to be dominated by internal motions), and $U$ is
the internal energy. For a spherically symmetric, constant density cloud, the
condition for virial stability is a quadratic one between velocity dispersion
and gas surface density, namely, $\sigma^{2}=(3/5)\pi GR_{cloud}\Sigma_{gas}$.
Figure 11a and Figure 11b show the line of sight velocity dispersion of the
gas as a function of the gas surface density for the Taffy bridge clouds, for
two values of assumed XCO. As in Figure 9, the color coding for the symbols
signifies the degree to which the observed molecular gas is exhibiting star
formation. The grey band in Figure 11 shows the range of loci of clouds in
virial equilibrium for diameters of 85-100 pc (Rcloud = 42.5-50 pc). Also
shown is the average relationship found by Leroy et al. (2016) for clouds on a
60pc scale for a sample of nearby normal galaxies from PHANGS, derived with
$X_{\rm CO,20}$. This latter relationship is flatter in this representation
than the quadratic one expected from the virial theorem. Thus the typical
clouds in normal galaxies show a slower growth888Leroy et al. (2016) found a
relationship of the form $log_{10}(\sigma)=0.31\leavevmode\nobreak\
\times\leavevmode\nobreak\ log_{10}(\Sigma_{gas,50})+0.74$, where $\sigma$ is
in km s-1, and gas surface density $\Sigma_{gas,50}$ is in units of 50 $\times
M_{\odot}pc^{-2}$.  in velocity dispersion $\sigma$ with increasing gas
surface density $\Sigma_{gas}$, suggesting that the gas surface density is not
the only factor at play in controlling properties in late-type disk galaxies
(see Leroy et al. 2016; Meidt et al. 2018, 2020).

We notice three aspects which are worthy of discussion in Figure 11. Firstly,
the blue stars associated with the brightest Pa$\alpha$ emission in the X-HII
region lie along the inner edge of the distribution, showing systematically
lower values of $\sigma$ for a given surface density of gas compared with the
other points. Secondly, the slope of the distribution of Taffy bridge points
in the figure is slightly steeper than both the relationship for normal
galaxies, and that of virial equilbrium relationship. A similar steeper trend
was found for gas clouds in the Antennae interacting system (Leroy et al.,
2016, not shown here) which shares some similarities with the Taffy. We will
discuss a comparison between the Antennae system and Taffy in § 6.1. Finally,
all the points in Figure 11a, and many of the points in Figure 11b lie on the
upper side of the virial equilibrium line, suggesting the clouds are not bound
by self-gravity. Indeed several observed profiles have high velocity
dispersion with FWHM (2.36$\sigma$) $\geq$ 100 km/s on the scale of 85 pc. At
XCO = 1/2 XCO,20, those points that lie on the bounded side of that plot tend
to have the strongest Pa$\alpha$, perhaps suggesting that those clouds have
overcome gravity to form stars.

Figure 12: The balance between bridge cloud internal motions $\sigma^{2}/R$
versus molecular surface density for the example case of $X_{\rm CO}=1/2X_{\rm
CO,20}$ and diameter = 85 pc (R = 42.5 pc). The colored symbols for the points
are the same as Figure 11. Model lines for equilibrium with an external
overpressure Pe from a confining medium are shown for $10^{6}<P_{e}/k<10^{9}$
in cgs units (from Field et al. 2011). The black line and red dotted lines
shows the condition for pure virial equilibrium for centrally concentrated and
constant-density clouds, respectively. The data for the protocluster candidate
cloud in the Antennae “overlap region”, sometimes called the “Firecracker”
(Johnson et al., 2015; Whitmore et al., 2014), is shown as a large yellow
filled circle, along with darker grey and grey-green filled points in the
lower left corner, for Galactic and nearby galaxies clouds respectively,
(adapted from Figure 6 of Johnson et al. 2015).

To explore what magnitude the overpressure would need to be to contain the
Taffy bridge gas, we show in Figure 12 another common representation of the
cloud kinematics (a logarithmic plot of $\sigma^{2}/R_{cloud}$ versus gas
surface density) for the case of denser clouds (X${}_{\rm CO}=1/2X_{\rm
CO,20}$). The virial stability condition is shown for cloud diameter 85 pc
(see caption). Also plotted are model curves for various over-pressures that
would be required to provide stability to the clouds through an external
medium999Increasing the assumed size of the clouds (to say 100 pc) would cause
the points in Figure 12 to move down and to the left by a relatively small
amount. More extended clouds would require less external pressure to bring
them into equilibrium.. Adapted from Figure 6 of Johnson et al. (2015), we
also show the distribution of clouds found in nearby galaxies and Galactic
clouds in the lower left corner of the plot. Those clouds tend to follow the
line of virial stability within the observational scatter. The blue filled
circles, which represent the hundreds of non-starforming Taffy bridge clouds,
would require very high pressures, as high as P/k of 106-8 (in cgs units), to
stabilize or overpressure them. As discussed by Field et al. (2011) and
Johnson et al. (2015) external pressures with Pe/k $>$ 105-6 are rarely found
in normal galaxies.

Could the direct high-speed collision of the Taffy galaxies create special
conditions that might allow for regions of high gas overpressure in an
otherwise relatively low-density turbulent medium? Silk (2019) has speculated
that high-speed cloud-cloud collisions, produced in galaxies collisions, can,
on the one hand, significantly suppress star formation in lower density
regions because of supersonic turbulence and shear in the gas, but on the
other hand, create rare over-pressured conditions (see Padoan et al. 2012).
Here, dense self-gravitating clouds become over-pressured, collapsing from a
slab-like compressed gas structure to form large numbers of “proto globular-
like” star clusters with a relatively high star formation efficiency (see also
Madau et al. 2020). Regions of high density in a supersonic turbulent field is
a natural result of the intermittency of turbulence, where highly non-linear
behavior can occur for short periods of time within the turbulent system (e.
g. Anselmet et al. 2001; Hily-Blant & Falgarone 2009). Since intermittency
occurs on scales close to the scale at which most of the dissipation is
occurring (probably pc-scales, see Guillard et al. 2009), it is not clear how
such high pressure regions within the multi-phase medium might affect the
collapse of gas on the larger scales needed to form super star clusters.

High overpressure from turbulence was suggested as an explanation for the
dense CO emitting clouds found within a larger supergiant molecular cloud
SGMC2 in the Antennae colliding galaxy system (Wilson et al., 2000), and
discussed in detail by Johnson et al. (2015). This cloud (dubbed the
“Firecracker”) would require even higher external pressure than the Taffy
clouds (see Figure 12 ).

###  5.4 Star formation in the X-HII region

Although the majority of the gas in the Taffy bridge is quite deficient in
star formation, the X-HII region is an exception. In Table 2 we quantify the
overall star formation properties of the X-HII region using a variety of
different measures, ranging from the Pa$\alpha$, H$\alpha$ and radio emission
measures, as well as infrared measures of star formation, or a combination of
both. Column 1 gives the observational method used, the radius (Column 2 and
3) of the circular aperture used in arcsec and area of the aperture in kpc2
(assuming D = 62 Mpc). Columns 4, 5 and 6 give the flux density, the measured
line fluxes integrated over the apertures for Pa$\alpha$ (this work) and
H$\alpha$ (from Joshi et al. 2019) and the total luminosity (for the IR this
is $\nu F_{\nu}$). Column 7 and 8 give the estimated SFR (before and after
extinction correction, where appropriate) based on methods shown in the table
notes. Column 9 presented the equivalent SFR surface density. We also estimate
the Pa$\alpha$ line and 6cm radio fluxes from the main extent of the X-HII
region and the two hotspots (Hotspot N and S) within the X-HII region
described earlier (Figure 7). These are given the designation ALL, N and S
respectively. The notes to the table give the method and reference source for
the various SFR measures used.

Table 2: Observed and derived properties of the Taffy extragalactic HII (X-HII) region Obs/Method | Rap | Area | Flux Density | Line Flux | LOG(L) | SFR | SFRcor | LOG($\Sigma_{SFR}$)
---|---|---|---|---|---|---|---|---
| ($arcsec$) | ($kpc^{2}$) | ($mJy$) | ($erg\leavevmode\nobreak\ s^{-1}cm^{-2}$ | ($erg\leavevmode\nobreak\ s^{-1}$) | ($M_{\odot}yr^{-1}$) | ($M_{\odot}yr^{-1}$) | ($M_{\odot}yr^{-1}kpc^{-2}$)
|  |  |  |  $\times 10^{-15}$) |  |  |  |
[1] | [2] | [3] | [4] | [5] | [6] | [7] | [8] | [9]
IRAC3.6$\mu$m | 4.8 | 6.4 | 0.38$\pm$0.03 | — | 41.18$\pm$0.04bbLuminosity $\nu$F($\nu)$, and assuming a distance to Taffy of 62 Mpc. We used $\nu$ = 8.3, 6.7, 5.2, 3.8 and 1.25 $\times$ 1013 Hz for the IRAC band 1, 2, 3 and 4 and MIPS 24$\mu$m respectively. | — | — | —
IRAC4.5$\mu$m | 4.8 | 6.4 | 0.32$\pm$0.03 | — | 40.98$\pm$0.04bbLuminosity $\nu$F($\nu)$, and assuming a distance to Taffy of 62 Mpc. We used $\nu$ = 8.3, 6.7, 5.2, 3.8 and 1.25 $\times$ 1013 Hz for the IRAC band 1, 2, 3 and 4 and MIPS 24$\mu$m respectively. | — | — | —
IRAC 5.8$\mu$m | 4.8 | 6.4 | 1.46$\pm$0.15 | — | 41.54$\pm$0.04bbLuminosity $\nu$F($\nu)$, and assuming a distance to Taffy of 62 Mpc. We used $\nu$ = 8.3, 6.7, 5.2, 3.8 and 1.25 $\times$ 1013 Hz for the IRAC band 1, 2, 3 and 4 and MIPS 24$\mu$m respectively. | — | — | —
IRAC 8$\mu$m | 4.8 | 6.4 | 4.33$\pm$0.43 | — | 41.86$\pm$0.04bbLuminosity $\nu$F($\nu)$, and assuming a distance to Taffy of 62 Mpc. We used $\nu$ = 8.3, 6.7, 5.2, 3.8 and 1.25 $\times$ 1013 Hz for the IRAC band 1, 2, 3 and 4 and MIPS 24$\mu$m respectively. | — | — | —
MIPS 24$\mu$m | 4.8 | 6.4 | 6.1$\pm$1.9 | — | 41.55$\pm$0.13 | 0.08$\pm$0.04ccAssuming a simple monochromatic relationship of Relaño et al. (2007); Calzetti et al. (2010), SFR${}_{24\micron}[M_{\odot}yr^{-1}]=5.66\times 10^{-36}$ $\times L_{\nu}(24\micron)^{0.82}$, where $L_{\nu}(24\micron)$ is in erg s-1. For the star formation surface density we assume the 24$\mu$m flux is emitted from the same area as the the P$\alpha$ flux, i.e. from an area of 0.69 kpc2.  | — | -0.94ccAssuming a simple monochromatic relationship of Relaño et al. (2007); Calzetti et al. (2010), SFR${}_{24\micron}[M_{\odot}yr^{-1}]=5.66\times 10^{-36}$ $\times L_{\nu}(24\micron)^{0.82}$, where $L_{\nu}(24\micron)$ is in erg s-1. For the star formation surface density we assume the 24$\mu$m flux is emitted from the same area as the the P$\alpha$ flux, i.e. from an area of 0.69 kpc2.
NIC 3 Pa$\alpha$ | 1.56 | 0.69 | — | 9.1$\pm$1.4 | 39.6$\pm$0.06 | 0.20 $\pm$0.03 | 0.24 $\pm$0.04iiassumes 0.2 mag of extinction at Pa$\alpha$ | -0.5
(ALL) |  |  |  |  |  |  |  |
Pa$\alpha$sf | 1.56 | 0.69 | — | 5.0$\pm$0.29eeassume both Pa$\alpha$ and H$\alpha$ fluxes have $\sim$ 45$\%$ contribution from shocks, so Pa$\alpha$sf (star formation) = 0.55 x Pa$\alpha$tot (Joshi et al., 2019) | 39.36$\pm$0.06eeassume both Pa$\alpha$ and H$\alpha$ fluxes have $\sim$ 45$\%$ contribution from shocks, so Pa$\alpha$sf (star formation) = 0.55 x Pa$\alpha$tot (Joshi et al., 2019) | 0.11$\pm$0.03d,ed,efootnotemark:  | 0.13$\pm$0.03i,d,ei,d,efootnotemark:  | -0.72
(ALL)eeassume both Pa$\alpha$ and H$\alpha$ fluxes have $\sim$ 45$\%$ contribution from shocks, so Pa$\alpha$sf (star formation) = 0.55 x Pa$\alpha$tot (Joshi et al., 2019) |  |  |  |  |  |  |  |
Pa$\alpha$ | 0.56 | 0.09 | — | 2.1$\pm$0.2 | 38.97$\pm$0.06 | 0.04$\pm$0.01 | 0.05iiassumes 0.2 mag of extinction at Pa$\alpha$ | -0.23
(Hotspot N) |  |  |  |  |  |  |  |
Pa$\alpha$ | 0.56 | 0.09 | — | 3.0$\pm$0.45 | 39.14$\pm$0.06 | 0.07$\pm$0.02 | 0.08iiassumes 0.2 mag of extinction at Pa$\alpha$ | -0.1
(Hotspot S) |  |  |  |  |  |  |  |
Pa$\alpha$+24$\mu$m | 1.56ffassume SFR (measured $H\alpha+24\micron)=5.3\times 10^{-42}[L(H\alpha_{sf})+0.031L(24\micron)]$, and that the 24$\mu$m flux come from the same area as the H$\alpha$ emission; (Calzetti et al., 2007). | 0.69ffassume SFR (measured $H\alpha+24\micron)=5.3\times 10^{-42}[L(H\alpha_{sf})+0.031L(24\micron)]$, and that the 24$\mu$m flux come from the same area as the H$\alpha$ emission; (Calzetti et al., 2007). | — | — | — | 0.15$\pm$0.05dfdffootnotemark:  | — | -0.66d,fd,ffootnotemark:
GCMS H$\alpha$ | 4.8 | 6.5 | — | 16.6$\pm$3.3 | 39.88$\pm$0.08 | — | — | —
(ALL) |  |  |  |  |  |  |  |
H$\alpha$sfeeassume both Pa$\alpha$ and H$\alpha$ fluxes have $\sim$ 45$\%$ contribution from shocks, so Pa$\alpha$sf (star formation) = 0.55 x Pa$\alpha$tot (Joshi et al., 2019) | 4.8 | 6.5 | — | 8.3$\pm$1.6eeassume both Pa$\alpha$ and H$\alpha$ fluxes have $\sim$ 45$\%$ contribution from shocks, so Pa$\alpha$sf (star formation) = 0.55 x Pa$\alpha$tot (Joshi et al., 2019) | 39.58$\pm$0.08eeassume both Pa$\alpha$ and H$\alpha$ fluxes have $\sim$ 45$\%$ contribution from shocks, so Pa$\alpha$sf (star formation) = 0.55 x Pa$\alpha$tot (Joshi et al., 2019) | 0.14$\pm$0.03 | 0.50$\pm$0.1 aaSFRcor: extinction corrected for Av = 1.5 mag based on the Balmer decrement of Joshi et al. (2019). | -1.11aaSFRcor: extinction corrected for Av = 1.5 mag based on the Balmer decrement of Joshi et al. (2019).
H$\alpha$sf+24$\mu$m | 4.8 | 6.5 | — | — | — | 0.09$\pm$0.02 ffassume SFR (measured $H\alpha+24\micron)=5.3\times 10^{-42}[L(H\alpha_{sf})+0.031L(24\micron)]$, and that the 24$\mu$m flux come from the same area as the H$\alpha$ emission; (Calzetti et al., 2007). | — | -1.86
Radio(6GHz) | 1.56 | 0.69 | 0.33$\pm$0.03 | — | [1.51$\pm$0.15 | 0.25$\pm$0.04ggfrom VLA 6 GHz flux densities and luminosities from Appleton et al. (in prep.) assuming the sum of thermal (T=104 K), non-thermal 6GHz contributions from Murphy et al. (2011), and calculated spectral index of $\gamma$ = -0.76 (this work) within a common area of 2.3 arcsec2 area.  | — | -0.41
(ALL) |  |  |  |  |  $\times$ 10${}^{20}W$Hz-1]ggfrom VLA 6 GHz flux densities and luminosities from Appleton et al. (in prep.) assuming the sum of thermal (T=104 K), non-thermal 6GHz contributions from Murphy et al. (2011), and calculated spectral index of $\gamma$ = -0.76 (this work) within a common area of 2.3 arcsec2 area.  |  |  | .
Radio(6GHz) | 0.56 | 0.08 | 0.052$\pm$0.005 | — | [2.38$\pm$0.24 | 0.040$\pm$0.002ggfrom VLA 6 GHz flux densities and luminosities from Appleton et al. (in prep.) assuming the sum of thermal (T=104 K), non-thermal 6GHz contributions from Murphy et al. (2011), and calculated spectral index of $\gamma$ = -0.76 (this work) within a common area of 2.3 arcsec2 area.  | — | -0.32
(Hotspot N) |  |  |  |  |  $\times$ 10${}^{19}W$Hz-1]ggfrom VLA 6 GHz flux densities and luminosities from Appleton et al. (in prep.) assuming the sum of thermal (T=104 K), non-thermal 6GHz contributions from Murphy et al. (2011), and calculated spectral index of $\gamma$ = -0.76 (this work) within a common area of 2.3 arcsec2 area.  |  |  | .
Radio(6GHz) | 0.56 | 0.09 | 0.086$\pm$0.009 | — | [3.94$\pm$0.40 | 0.07$\pm$0.01ggfrom VLA 6 GHz flux densities and luminosities from Appleton et al. (in prep.) assuming the sum of thermal (T=104 K), non-thermal 6GHz contributions from Murphy et al. (2011), and calculated spectral index of $\gamma$ = -0.76 (this work) within a common area of 2.3 arcsec2 area.  | — | -0.11
(Hotspot S) |  |  |  |  |  $\times$ 10${}^{19}W$Hz-1]ggfrom VLA 6 GHz flux densities and luminosities from Appleton et al. (in prep.) assuming the sum of thermal (T=104 K), non-thermal 6GHz contributions from Murphy et al. (2011), and calculated spectral index of $\gamma$ = -0.76 (this work) within a common area of 2.3 arcsec2 area.  |  |  | .
Radio(1.4GHz) | 1.56 | 0.69 | 1.2$\pm$0.1 | — | [5.17$\pm$0.50 | 0.29$\pm$0.03hhfrom VLA 1.48 GHz flux density and luminosity of Appleton et al (in prep.), and assuming the SFR prescription of (Condon, 1992) and non-thermal fraction of flux = 0.8. The source is very extended at 20cm compared with the restoring beam = 1.27 x. 1.1 arcsec2. | — | -0.38
(ALL) |  |  |  |  |  $\times$ 10${}^{20}W$Hz-1]hhfrom VLA 1.48 GHz flux density and luminosity of Appleton et al (in prep.), and assuming the SFR prescription of (Condon, 1992) and non-thermal fraction of flux = 0.8. The source is very extended at 20cm compared with the restoring beam = 1.27 x. 1.1 arcsec2. |  |  | .

The SFRs estimated in Table 2 provide a relatively consistent picture
considering the different methods used. Joshi et al. (2019) showed that the
H$\alpha$ line emission is contaminated by up to 45$\%$ by shocked gas in the
bridge, and so the line fluxes for Pa$\alpha$ have been reduced accordingly
before calculating the SFRs. For the entire X-HII region, we can compare the
shock corrected Pa$\alpha$ SFR, Pa$\alpha$ (ALL) = 0.13 $M_{\odot}$ yr-1, with
that derived from the extinction independent 6 GHz radio continuum,
Radio(6GHz) (ALL) = 0.25 $M_{\odot}$ yr-1. A low SFR is supported using a pure
MIPS 24$\mu$m flux calculation101010We note that the background-subtracted
24$\mu$m flux associated with the X-HII region is uncertain because of
contamination by the disk of UGC 12915. Uncertainties in our estimated
contamination are included in the uncertainty in Table 2. of 0.08 $M_{\odot}$
yr-1. Calzetti et al. (2007) provide a composite method of calculating the SFR
using the 24$\mu$m flux in combination with the H$\alpha$ emission uncorrected
for extinction. Using either the observed Pa$\alpha$ flux as a proxy for the
H$\alpha$ flux, or the observed H$\alpha$ flux in combination with 24$\mu$m,
yields 0.15 and 0.09 $M_{\odot}$ yr-1 respectively. These values are in
contrast to the extinction corrected H$\alpha$ flux measurements which yield
H$\alpha_{sf}$ = 0.5 $M_{\odot}$ yr-1. The high value for H$\alpha_{sf}$ may
imply that Joshi et al. (2019) overestimated the optical extinction at
H$\alpha$. We conclude that the most likely SFR for the main Pa$\alpha$
emitting area of the X-HII region is in the range 0.10-0.25
$M_{\odot}\leavevmode\nobreak\ yr^{-1}$.

We estimate that the molecular gas mass associated with the entire X-HII
region (4.8 arcsec2) based on our CO observations is 1.29$\pm$0.01 $\times$
108 [$X_{\rm CO}/X_{\rm CO,20}$] $M_{\odot}$. Formally, the depletion time for
the molecular gas, $M_{mol}$/SFR, is therefore 645 Myrs and 160 Myrs,
respectively for $X_{\rm CO}=X_{\rm CO,20}$, and $X_{\rm CO}=1/4X_{\rm
CO,20}$. However this depletion time is likely much longer because the star
formation is quite clumpy and inhomogeneous. The KS-diagrams of Figure 9a+b,
show that where many of the brighter SF regions lie close to the the average
KS relationship, the implied depletion times are $\sim$ 1Gyr.

The Pa$\alpha$ hotspots in the X-HII region could be examples of the formation
of massive proto-clusters conceptually similar to fast gas-on-gas collisions
between dwarf galaxies postulated as one formation mechanism for Ultra Diffuse
Galaxies (Silk 2019; UDFs). Table 2 provides estimates of the SFR for the
north (N) and south (S) hotspots of 0.05 $M_{\odot}\leavevmode\nobreak\
yr^{-1}$ and 0.08 $M_{\odot}\leavevmode\nobreak\ yr^{-1}$ respectively, based
on Pa$\alpha$. The radio continuum provides approximately similar values. Over
the next 10 Myrs, 10${}^{6}M_{\odot}$ of stars could form, making them future
young massive star cluster candidates.

##  6 Other potentially similar environments to the Taffy bridge gas

###  6.1 Comparison with Antennae “overlap” region

It is worthwhile contrasting the properties of the Taffy bridge with another
well studied colliding galaxy system, the Antennae galaxy pair, NGC4038/9. The
Antennae is known to be in the starburst phase of a major merger (Whitmore et
al., 1999; Schweizer et al., 2008). Unlike the Taffy system, which has
experienced a head-on counter-rotating collision in the past (25-35 Myr ago),
the Antennae is still in the compressive “contact“ stage of a relatively slow
(100 Myr) prograde disk-disk merger (Renaud et al., 2018). In global
properties, the Antennae and Taffy have very similar total far-infrared (FIR)
luminosities and total molecular gas masses. The Antennae has a FIR
luminosity, LFIR= 5.6 $\times$ 10${}^{10}L_{\odot}$ (Brandl et al., 2009) and
total molecular mass $1.1\leavevmode\nobreak\ \times\leavevmode\nobreak\
10^{10}M_{\odot}$ (Zhu et al., 2003), whereas equivalent values for the Taffy
are LFIR = 6.5 $\times\leavevmode\nobreak\ 10^{10}L_{\odot}$ (Sanders et al.,
2003)) and a total molecular mass $M_{\rm mol}=0.97\leavevmode\nobreak\
\times\leavevmode\nobreak\ 10^{10}M_{\odot}$ (Zhu et al., 2007). The “overlap”
region of the Antennae (which lies between the two component galaxies) also
shares a number of similarities with the Taffy bridge. They both include
significant quantities of molecular gas distributed in narrow filaments.
Bearing in mind the uncertainties in the XCO factor for both galaxies, the
Taffy bridge has roughly twice as much molecular gas as the “overlap” region.
The Taffy bridge contains $\sim 1.3\leavevmode\nobreak\
\times\leavevmode\nobreak\ 10^{9}M_{\odot}$ (for X${}_{\rm
CO}=1/5\leavevmode\nobreak\ X_{\rm CO,20}$; Zhu et al. 2007), compared with
$\sim 0.5\leavevmode\nobreak\ \times\leavevmode\nobreak\ 10^{9}M_{\odot}$ for
the Antennae ”overlap“ region (Stanford et al. 1990; Wilson et al. 2000, if we
assume a common XCO factor. Unlike the Taffy bridge, which we have shown is
largely devoid of star formation, the Antennae overlap region is bursting with
activity. The total star formation rate in the overlap regions has been
measured to be 5 $M_{\odot}$ yr-1 (Stanford et al., 1990) compared with
0.1-0.25 $M_{\odot}$ yr-1 for the Taffy bridge (this paper). In the Antennae
overlap region, the star formation activity is concentrated in a handful of
very massive Super Giant Molecular Complexes (Wilson et al., 2000). With the
one exception of the “Firecracker” discussed earlier, the majority of these
giant molecular complexes in the Antennae dominate the star formation output
of the overlap region (Mirabel et al., 1998; Whitmore & Schweizer, 1995;
Whitmore et al., 2010; Johnson et al., 2015). In the Taffy, the molecular
clouds are less massive, and only the X-HII region compares in molecular mass
(108 $M_{\odot}$) with the lower end of the mass distribution of the Antennae
SGMCs. Another difference is that, unlike the Taffy filaments, which contain
clouds with locally high velocity dispersion (40-100 km s-1), the Antennae CO
filaments exhibit unusually low velocity dispersion ($\leq$ 10 km s-1;
Whitmore et al. 2014). Higher CO velocity dispersion are reported for the
Antennae system in general (Leroy et al., 2016), which likely correlate with
increased star formation activity. This seems the opposite of the Taffy, where
the X-HII region exhibits the lowest velocity dispersion in the CO clouds,
whereas the filaments and star formation-deficient clouds have the highest
velocity dispersion. In summary, except for some regions of high star
formation rate in the Taffy X-HII region, the filaments and clumps in the
Taffy bridge are strongly suppressed in star formation compared with the
Antennae.

###  6.2 Comparison with molecular gas in ram-pressure stripped systems

Another kind of environment that might have some similarities with the Taffy
bridge are those involving the ram-pressure stripping of gas from cluster
galaxies, where stripped molecular clouds find themselves in the intergalactic
medium. There is now significant evidence that the ISM of late-type cluster
galaxies can be stripped (Gunn & Gott, 1972) by their interaction with a hot
gaseous cluster halo, including the formation of tails revealed in HI,
H$\alpha$ emission, X-rays and warm and cold molecular hydrogen (e. g. Gavazzi
et al. 2001; Wang et al. 2004; Oosterloo & van Gorkom 2005; Sun et al. 2007;
Yagi et al. 2007; Jáchym et al. 2014). Although the physical mechanism from
stripping gas into the intergalactic medium (see Schulz & Struck 2001) is
rather different from the case of a major head-on collision, as in the Taffy
case, there are some potential similarities. For example, Sivanandam et al.
(2010) found evidence for warm molecular hydrogen in the stripped tail of ESO
137-001 that was likely shock-heated by turbulent interactions with its host
ICM of Abell cluster 3627. Cold molecular gas has also been found in some
likely ram-pressure stripped galaxies, primarily through single dish
observations of low spatial resolution (e. g. Jáchym et al. 2014, 2017; Lee et
al. 2017). Only recently have high resolution observations become available
(Jáchym et al., 2019) including those studied by the GASP survey111111GASP =
Gas Stripping Phenomena with MUSE (Poggianti et al., 2017).. In these cases,
these intergalactic clouds, either stripped from their host galaxies or formed
in situ through turbulent compression, can be compared with the Taffy bridge.

One particular gas-stripped system, JW 100, is a member of a class of similar
“Jellyfish” galaxies (Ebeling et al., 2014), and was observed with ALMA in the
CO (1-0) and (2-1) transitions (Moretti et al., 2018, 2020). This massive
cluster galaxy (with a diameter of at least 45 kpc) exhibits a clumpy tail of
molecular hydrogen extending 35 kpc away from its disk. Approximately
$7\leavevmode\nobreak\ \times\leavevmode\nobreak\ 10^{9}M_{\odot}$ of
molecular gas is present in the tail, which is comparable with the Taffy
bridge, representing about 30$\%$ of the gas in the galaxy disk. Like the
Taffy, several clumps in the tail exhibit high velocity dispersion ( 70-80 km
s-1), although much of the gas exhibits lower linewidths. Although the
mechanism for stripping the gas in such tails may be markedly different from
the Taffy (Pedrini et al., 2022), it is possible that gas may be formed in
shocks or eddies in the wake of the main shocked region around the galaxy. Is
there evidence for star formation suppression here too?

In Figure 13 we show examples taken from Moretti et al. (2020), of clouds in
the stripped tail of the JW100 (blue circles) overplotted on a partially
greyed-out version of one characteristic KS-relationship plots of the Taffy
from this paper (Figure 10b). As is discussed in more detail by Moretti et al.
(2020), many of the stripped clouds fall below the standard KS relationship.
Thus, like the majority of non-detected Pa$\alpha$ upper limits in Taffy, the
stripped clouds seem to show very low star formation rates, as measured by
optical spectroscopy. One difference between those stripped clouds and Taffy
is that the former appear to have significantly lower overall H2 surface
densities compared with the Taffy bridge, even if (as was assumed by the
authors) XCO has a galactic value. However, one should caution that the ALMA
data discussed by Moretti et al. (2020) was taken with 1 arcsec resolution,
which corresponds to 1kpc at the distance of JW100. Thus the surface densities
of molecular gas could well be much higher than observed, if the system had
been observed at the same linear resolution as the Taffy. Nevertheless, it is
clear that there are some similarities between the properties of the gas in
the stripped clouds in the JW 100 and those suppressed star formation clouds
in the Taffy bridge.

The careful comparison of the multi-wavelength observations of JW100 by
Poggianti et al. (2019) further strengthens those similarities. Like the Taffy
X-HII region, an ULX source was found associated with one of the brighter HII
regions. These authors also showed that although parts of the stripped gas
show evidence of star formation, there are regions of the tail that show
LINER-type optical emission (similar to regions of the Taffy bridge Joshi et
al. 2019), and X-ray emission not consistent with star formation. A variety of
possible heating mechanisms for such systems, including strong turbulent
mixing, plasma interactions, thermal conduction, and shocks have been put
forward as possible explanations (Poggianti et al., 2019; Campitiello et al.,
2021; Pedrini et al., 2022).

Figure 13: A comparison between the surface formation rate density and
molecular hydrogen surface density for clouds in the ram-pressure stripped
tail of the jellyfish galaxy JW100 (blue filled circles Moretti et al. 2020,
where XCO was assumed to be Galactic), and a partially greyed-out version of
Figure 10b of the current paper (symbols are the same as that figure). Like
the majority of the Taffy clouds, which have upper limits to star formation,
the JW100 clouds also show significant star formation suppression.

##  7 Origin of the Star Formation Suppression in the Bridge

The majority of the clouds in the Taffy bridge are likely unbound if self-
gravity is the only restoring force for a wide range of reasonable values of
assumed XCO. Such clouds would evaporate on a crossing time ($2R/\sigma$),
which is typically 2-8 Myr (see Figure 14) for the average Taffy bridge cloud.
This is much shorter then the nominal time since the Taffy galaxies collided
of 25-35 Myrs (Condon et al., 1993). This short timescale seems to rule out
the possibility that the clouds in the bridge were originally normal molecular
clouds confined to their host disks before the collision. In such a simple
picture, such clouds would suddenly no longer feel the gravitational forces of
the original disk, and maybe expected to freely expand when ejected into
intergalactic space. Although it is possible that a sub-set of clouds (the
long tail of the distribution in Figure 14) may have crossing times long
enough to have survived without an external medium since the collision (we
cannot resolve the linewidths of clouds with FWHM $<$ 15 km s-1), it is clear
that the majority of the gas, including those associated with the X-HII
region, have crossing times which suggest they are relatively recently formed.
Furthermore, models of the head-on collision like the Taffy system show
significant ionization of the ISMs in both galaxies at and shortly after the
collision suggesting that only a relatively small number of clouds punch
through unscathed (Yeager & Struck, 2020a, b).

Figure 14: Cloud crossing times in the bridge in Myr for an assumed diameter
of 85 pc. If the clouds are smaller the crossing times will be even shorter.
The clouds devoid of star formation (solid blue) have generally shorter
crossing times than those containing star formation (solid red). The other
samples are the same as those described in Figure 9.

To understand why the CO clouds are not virialized, we are reminded that the
Taffy bridge CO clouds are embedded in an highly turbulent multi-phase medium.
There is evidence for the existence of copious quantities of shock-heated warm
H2 through direct mid-IR emission lines, and indirectly through the detection
of [CI] and [CII] emission (Peterson et al., 2012, 2018). Additionally, there
is evidence for shock-excited atomic emission line gas in the bridge (Joshi et
al., 2019). Thus the CO-emitting clouds are likely immersed in a highly
turbulent medium, similar to the intergalactic shocked filament in Stephan’s
Quintet (Appleton et al., 2017; Guillard et al., 2021). Vollmer et al. (2021)
has discussed this in the context of their dynamical model of the Taffy system
which was tuned to provide a good match with the lower-resolution PdBI CO
observations of the Taffy. In that paper, the authors use their model to show
that turbulence driven by cloud-cloud collisions on large-scales in the bridge
would lead to clouds which are not virialized. Our observation, which push
down by an order of magnitude in scale compared with the PdBI observations,
fully support this conclusion.

However the actual details of the mechanism by which cold CO clouds are
dynamically heated are likely to be quite complex. For example, the cascade of
energy from the large driving scales of the collision, and subsequent
continued high Mach-number cloud-cloud collisions can only be properly modeled
with multi-phase model which takes into account atomic and molecular cooling
phases, phase transitions, shocks, turbulence and gas phase-mixing in
different media (Guillard et al., 2009). For example, in the analogous
Stephan’s Quintet shocked filament, recent observations of Ly$\alpha$ emission
(Guillard et al., 2021) combined with previous observations of molecular
hydrogen and X-rays (Appleton et al., 2017; O’Sullivan et al., 2009) show
strong evidence that the dissipation of kinetic energy involves all gas
phases. These are likely connected through a turbulent cascade involving
atomic and molecular magneto-hydrodynamic shocks (Lesaffre et al., 2020;
Durrive et al., 2021), and turbulent mixing layers (Slavin et al., 1993). The
models of Yeager & Struck (2020a, b) are a step in the right direction in
attempting to treat multi-phase cooling of the gas in Taffy-like bridge
systems. Those models predict rapid evolution in the gas phases as the
collision progresses as shocks permeate the medium, with the appearance of
waves of cooling, and periods of re-excitation by cloud-cloud collisions. The
inclusion of mid-IR H2 and [CII] cooling lines is another important step
(Yeager et al, in preparation). These models, however, cannot yet capture the
turbulent mixing that must occur at the boundaries between the hot, warm and
cold components. Such models only partially deal with the possibility that the
cold molecular clouds are transient material that is constantly being recycled
through different thermal phases. In such cases, clouds must cool and collapse
rapidly to have a chance of forming stars.

The existence of high Mach-number gas collisions may stochastically generate
regions like the X-HII region (Yeager & Struck, 2020b). The Pa$\alpha$
hotspots in the X-HII region could be examples of the formation of massive
proto-clusters conceptually similar to fast gas-on-gas collisions between
dwarf galaxies postulated as one formation mechanism for Ultra Diffuse
Galaxies (Silk 2019; UDFs). The hotspots have locally high star formation
surface densities, and over the next 10 Myrs the extragalactic HII region
could form $\sim$ 1.3 $\times 10^{6}\leavevmode\nobreak\ M_{\odot}$ of stars,
making them future young massive star cluster candidates. Although in the case
of the Taffy, where the size of the colliding systems is larger than Silk
(2019) envisaged, high Mach-number gas collisions will likely continue to
generate regions like the X-HII region across the full extent of the bridge
over time (Yeager & Struck, 2020b).

An added complexity is the possible influence of magnetic fields. The Taffy
was discovered because of the detection of radio continuum between the
galaxies resulting from synchrotron radiation from cosmic rays spiraling in
magnetic fields (Condon et al., 1993). Lisenfeld & Völk (2010) explained the
radio emission as the result of cosmic rays accelerated in shocks produced in
situ in the Taffy bridge. The powerful emission from the lowest pure
rotational transitions of H2 in Stephan’s Quintet, which exhibits very similar
Mid-IR H2 spectra to the Taffy bridge, is most easily explained (Lesaffre et
al., 2013) by magnetic shocks which moderate the temperature of the shock into
the temperature range where the 0-0S(0) and 0-0S(1) line dominate (100 $<$ T
$<$ 300 K). New deep VLA radio continuum observations with much higher
resolution than the original Condon et al. observation have been recently made
(Appleton, in preparation) which will provide a more complete picture of the
distribution of thermal and non-thermal emission in the Taffy bridge, and may
help determine the degree of importance of magnetic fields in the process of
energy dissipation and cloud dynamics.

##  8 Conclusions

This is the first of two papers describing ALMA observations of CO in the
Taffy galaxy system, UGC 12914/5. In this paper we present a large mosaic of
primary beam pointings made in the CO (2-1) transition at a spatial resolution
more than an order of magnitude greater than previous observation (0.24 x 0.18
arcsec2). The paper concentrates on the molecular bridge extending between the
galaxies. When compared with archival HST NICMOS Pa$\alpha$ and Spitzer
observations, as well as new VLA radio continuum, Palomar 5-m optical
spectroscopy and previously published Chandra X-ray observations, we conclude
the following:

  * •

The observations provide evidence that gas on scales of 60-100 pc in the Taffy
galaxies is significantly disturbed in the aftermath of the head-on collision
between the galaxies. The gas in UGC 12915 is particularly disturbed, with
streamers of likely tidal debris extending northwards far from the inner disk.
UGC 12914, although less obviously disturbed that its companion, exhibits a
singular powerful narrow structure of CO which follows the southern edge of
the inner disk into a hook-like arm.

  * •

The molecular gas bridge between UGC 12914/5 is resolved into a myriad of
narrow filaments (60-100pc $\times$ 1kpc) and clumpy structures extending
between the two galaxies which have some similarities with the filaments seen
in the “overlap” region of the Antennae galaxies, except that almost all of
them are devoid of star formation. By analyzing 234 regions in the bridge, we
compare the star formation rate surface density derived for Pa$\alpha$ HST
observations with the ALMA data at 85 pc resolution. We show that the majority
of the filaments are devoid of star formation, and fall significantly below
the Kennicutt-Schmidt relationship for normal galaxies, especially for the
numerous regions undetected in Pa$\alpha$. This statement is shown to be true
over a wide range of assumed extinction, XCO values and dust to gas ratios.
The result strongly supports the idea that, except in the X-HII region, the
gas in the bridge is experiencing strong star formation suppression, even at
the high surface densities probed by the new ALMA observations.

  * •

The kinematics of the majority of the bridge molecular gas shows unusually
high velocity dispersion on the scale of 85 pc, with some regions within the
filaments showing Gaussian line profiles with FWHM $\geq$ 100 km s-1. Gas
associated with the X-HII region tends to have lower velocity dispersion,
suggesting it is less turbulent there. Like clouds seen in the Antennae
galaxies, the Taffy clouds show a steeper slope in log($\sigma$) versus
log($\Sigma_{gas}$) than normal galaxies. In addition, most of the Taffy
bridge regions are gravitationally unbound for most reasonable values of
assumed $X_{\rm CO}\leavevmode\nobreak\ (1/4<(X_{\rm CO}/X_{\rm CO,20})<1/2$),
and, like the “Firecracker” region in the Antennae system, would require an
extremely high external pressure ($10^{6}<P_{e}/k<10^{8}$) to pressure-support
them. If they are not externally supported, many of the clouds would disperse
on a crossing timescale of $\sim$2-5 Myrs and are therefore likely transient.
Such clouds may be continuously created and destroyed in such a highly
turbulent medium.

  * •

Despite the highly turbulent medium, stars are able to sometimes form there,
albeit at a low rate. In the X-HII region we derive a star formation rate of
0.1-0.25 $M_{\odot}$ yr-1. Gas associated with this region exhibits high local
star formation rates, on or above the KS relationship for normal systems
depending on the assumed XCO value. These higher star formation rate densities
of 0.6-0.8 $M_{\odot}$ yr-1 kpc-2 in two regions (containing faint compact
radio sources and a ULX source) on scales of $<$ 100 pc might be evidence of
the formation of massive star clusters. We also present evidence of potential
rotational motions in the molecular gas and ionized gas associated with the
X-HII region.

The discovery that the majority of the dense molecular clouds in the bridge
are not able to form stars suggests that turbulence, driven by the collision
of the two Taffy galaxies more than 25-30 Myr previously, is still actively
suppressing star formation down the the scale of 60-100 pc across the majority
of the bridge. Recent gas-dynamical models of Taffy-like head-on collisions
support the idea that even after 30 Myr, cloud-cloud collisions are still
creating high Mach-number collisions in the bridge. Such collisions may also
explain the large quantities of warm molecular hydrogen discovered in the
bridge by the Spitzer IRS, as well as the existence of extensive shock-excited
ionized gas throughout the bridge, including strong H$\alpha$ emission at a
discrepant velocity near the X-HII region in the bridge. The process in which
large-scale driven turbulence caused by the galaxy collision can influence
dense cold molecular clouds on 60-100 pc scales is likely to involve many
different processes, including atomic and molecular magnetic shocks and
turbulent mixing of gas, similar to that suspected in the Stephan’s Quintet
system.

This paper made use of data from the following observatories/instruments: The
Atacama Large Millimeter/submillimeter Array (ALMA), the Hubble Space
Telescope (HST; NICMOS), the Chandra X-ray telescope S3 Advanced CCD Imaging
Spectrometer, the Very Large Array (VLA), the Spitzer Space telescope; IRAC
and MIPS and the Double-beam Spectrograph on the 5-m Hale telescope at the
Palomar Observatory. The paper made use of the following software: Community
Astronomy Software Applications (CASA; NRAO), Interactive Data Language (IDL;
L3Harris Geospatial Solutions), Python 3 open source software, PypeIt
(Prochaska et al., 2020), SAOImage DS9 (Chandra X-ray Science Center, HEASARC
and JWST Mission office at STSCI.)

This paper makes use of the following ALMA data: ADS-JAO.ALMA.2016.1.01037.S.
ALMA is a partnership of ESO (representing its member states), NSF (USA) and
NINS (Japan), together with NRC (Canada), MOST and ASIAA (Taiwan), and KASI
(Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA
Observatory is operated by ESO, AUI/NRAO and NAOJ. The DBSP spectra are based
on observations obtained at the Hale Telescope, Palomar Observatory as part of
a continuing collaboration between California Institute of Technology,
NASA/JPL, Yale University, and the National Astronomical Observatories of
China. This work also contains archival data obtained with the Spitzer Space
Telescope, which was operated by the Jet Propulsion Laboratory, California
Institute of Technology under a contract with NASA. The National Radio
Astronomy Observatory is a facility of the National Science Foundation
operated under cooperative agreement by Associated Universities, Inc.. This
research has made use of the NASA/IPAC Infrared Science Archive, which is
funded by the National Aeronautics and Space Administration and operated by
the California Institute of Technology. PA would like to acknowledge the NRAO
visitor support during a visit to NRAO-Charlottesville early in this project.
PA would also like to thank the Institut d’Astrophysique de Paris for support
as a Visiting Scientist in October 2019 in connection with this work. UL
acknowledges support from project PID2020-114414GB-100, financed by
MCIN/AEI/10.13039/501100011033, from projects P20_00334 and FQM108, financed
by the Junta de Andalucia and from FEDER/Junta de Andalucía-Consejería de
Transformación Económica, Industria, Conocimiento y Universidades/Proyecto
A-FQM-510-UGR20. The authors wish to thank an anonymous referee for many
important suggestions for improving this paper.

##  Appendix A Extracted Regions

In this Appendix we provide sample Tables (in machine readable format) of the
regions extracted in the bridge and the measured and derived properties. The
full contents of the tables will be made available online. We provide here
only a sample of the full table. For the refereed version we include the full
tables. Column 1 provides the Region which corresponds to the regions marked
in Figures 16 and 17, and were chosen to be representative regions in the
bridge structures. Column 2 and 3 are the RA (J2000) and Dec (J2000) of the
regions. Column 4 and 5 are the best fitting systemic heliocentric velocity
(optical definition) and full-width half-maximum (FWHM), with uncertainties,
of the CO profiles after fitting with a Gaussian in km s-1. Column 6 is the
integrated flux of the CO (2-1) profile in Jy km s-1and uncertainties. Column
7 is a spectral quality flag based. Flags of unity are the best quality
spectra, and the table notes explain the other flags. Only Quality 1 flagged
data are used in the figures presented in the paper. Column 8 provides the
molecular mass surface density, in units of log10($M_{\odot}$ pc-2), for the
nominal case of XCO = 1/4 XCO,20. In the paper we consider potentially
realistic cases of XCO in the range 1/4 $<$ XCO,20 $<$ 1/2, where X${\rm
CO,20}$ is the standard Galactic value. Column 9 tabulates the surface density
of Pa$\alpha$ emission in units of log10(erg s-1 cm-2 kpc-2) for the Case 1
extinction example given in §5.1. Column 10 gives the star formation surface
density in units of log10($M_{\odot}$ yr-1 kpc-2) under the same assumptions.
Column 11 provides a Pa$\alpha$ detection flag (see table note).

Figure 15: Examples of spectra described in Table 3 and in Figure 16 and 17.
The designation refers to the entry in the Table and the blue lines show
Gaussian fits to the spectra. Narrow lines are seen in the region associated
with the X-HII region (e. g. D9 and D12). Table 3: Properties of the Extracted
Regions

Reg. | RA | Dec |  $V_{sys}$bbHeliocentric radial velocity using the optical definition. | FWHMccFWHM corrected for instrumental resolution. | $\Sigma S_{V}\Delta V$ | Specddassume CASE B recombination, f(H$\alpha$) = 2.86/0.332 x f(Pa$\alpha$), and the SFR - H$\alpha$ relations of Calzetti et al. (2007, 2010): SFR($[M_{\odot}yr^{-1})=5.3\times 10^{-42}L(H\alpha_{corr})$, where L(H$\alpha_{corr})$ is extinction corrected in erg s-1. | Log($\Sigma_{mol}$)eeMolecular surface density for nominal case of XCO = 1/4 XCO,20, see text. |  $\Sigma$(Pa$\alpha$)ffSurface flux and star formation rates estimated from the Pa$\alpha$ emission assuming minimal (Case 1) extinction (see §5.1).  | Log($\Sigma_{SFR})$ffSurface flux and star formation rates estimated from the Pa$\alpha$ emission assuming minimal (Case 1) extinction (see §5.1).  | Pa$\alpha$ggPa$\alpha$ detection flags; Flag = 0 indicates an upper limit, Flag = 1 indicates a detection, and Flag = -1 indicates that the region lies outside the Pa$\alpha$ NICMOS coverage area.
---|---|---|---|---|---|---|---|---|---|---
$\\#$ | (J2000) | (J2000) | CO (2-1) | CO (2-1) | CO (2-1) | Qual. | 1/4 XCO,20 | ($10^{-16}\times erg\leavevmode\nobreak\ s^{-1}$ | ($M_{\odot}yr^{-1}$ | Det.
| (hh:mm:ss) | (dd:mm:ss) | ($km\leavevmode\nobreak\ s^{-1}$) | ($km\leavevmode\nobreak\ s^{-1}$) | ($Jy$ km s-1) |  | ($M_{\odot}\leavevmode\nobreak\ pc^{-2}$) |  $cm^{-2}\leavevmode\nobreak\ {arcsec}^{-2}$) |  $kpc^{-2}$) | Flag
[1] | [2] | [3] | [4] | [5] | [6] | [7] | [8] | ([9] | [10] | [11]
A1a | 00:01:41.479 | +23:29:37.1 | 4380 | 96$\pm$14 | 0.29$\pm$0.06 | 1.0 | 2.46$\pm$0.08aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0.43aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | -1.75aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0
A1b | 00:01:41.479 | +23:29:37.1 | 4510 | 28$\pm$4 | 0.08$\pm$0.02 | 1.0 |  |  |  | 0
A2a | 00:01:41.466 | +23:29:37.0 | 4376 | 32$\pm$5 | 0.09$\pm$0.02 | 1.0 | 2.27$\pm$0.08aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0.43aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | -1.75aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0
A2b | 00:01:41.466 | +23:29:37.0 | 4519 | 28$\pm$3 | 0.14$\pm$0.02 | 1.0 |  |  |  | 0
A3a | 00:01:41.463 | +23:29:36.7 | 4375 | 36$\pm$7 | 0.09$\pm$0.02 | 1.0 | 2.23$\pm$0.11aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0.43aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | -1.75aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0
A3b | 00:01:41.463 | +23:29:36.7 | 4522 | 30$\pm$3 | 0.13$\pm$0.01 | 1.0 |  |  |  | 0
A4a | 00:01:41.448 | +23:29:36.4 | 4376 | 41$\pm$6 | 0.11$\pm$0.02 | 1.0 | 2.33$\pm$0.08aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0.43aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | -1.75aaRegions A1 to A4 have multiple CO components 1a, 1b, 2a, 2b.. etc. Their integrated Pa$\alpha$ flux and SFR properties are combined.  | 0
A4b | 00:01:41.448 | +23:29:36.4 | 4508 | 47$\pm$3 | 0.16$\pm$0.02 | 1.0 |  |  |  | 0
A5 | 00:01:41.451 | +23:29:36.1 | 4481 | 57.8$\pm$5 | 0.20$\pm$0.02 | 1.0 | 2.20$\pm$0.05 | 0.43 | -1.75 | 0
A6 | 00:01:41.423 | +23:29:36.5 | 4376 | 67$\pm$5 | 0.23$\pm$0.03 | 1.0 | 2.26$\pm$0.05 | 0.43 | -1.75 | 0
A7 | 00:01:41.424 | +23:29:36.3 | 4390 | 86$\pm$17 | 0.17$\pm$0.05 | 1.0 | 2.14$\pm$0.11 | 0.43 | -1.75 | 0
A8 | 00:01:41.429 | +23:29:35.7 | 4496 | 95.8$\pm$24 | 0.23$\pm$0.08 | 1.0 | 2.25$\pm$0.13 | 0.43 | -1.75 | 0
A9 | 00:01:41.445 | +23:29:35.6 | 4496 | 63.8$\pm$5 | 0.22$\pm$0.02 | 1.0 | 2.24$\pm$0.05 | 0.43 | -1.75 | 0
A10 | 00:01:41.435 | +23:29:35.3 | 4484 | 70.5$\pm$6 | 0.24$\pm$0.03 | 1.0 | 2.27$\pm$0.05 | 0.43 | -1.75 | 0
A11 | 00:01:41.393 | +23:29:35.2 | 4446 | 52$\pm$4 | 0.27$\pm$0.03 | 1.0 | 2.34$\pm$0.05 | 0.43 | -1.75 | 0
A12 | 00:01:41.377 | +23:29:35.2 | 4450 | 62$\pm$5 | 0.28$\pm$0.03 | 1.0 | 2.34$\pm$0.05 | 0.43 | -1.75 | 0
A13 | 00:01:41.396 | +23:29:35.0 | 4441 | 52$\pm$4 | 0.27$\pm$0.03 | 1.0 | 2.32$\pm$0.05 | 0.43 | -1.75 | 0
A14 | 00:01:41.380 | +23:29:35.0 | 4443 | 57$\pm$5 | 0.26$\pm$0.03 | 1.0 | 2.32$\pm$0.05 | 0.43 | -1.75 | 0
A15 | 00:01:41.365 | +23:29:34.2 | 4409 | 45$\pm$4 | 0.16$\pm$0.02 | 1.0 | 2.10$\pm$0.05 | 0.43 | -1.75 | 0
A16 | 00:01:41.446 | +23:29:33.9 | 4459 | 64$\pm$5 | 0.25$\pm$0.03 | 1.0 | 2.30$\pm$0.05 | 0.43 | -1.75 | 0
A17 | 00:01:41.463 | +23:29:34.1 | 4492 | 68$\pm$5 | 0.26$\pm$0.03 | 1.0 | 2.32$\pm$0.05 | 0.43 | -1.75 | 0
A18 | 00:01:41.342 | +23:29:33.7 | 4373 | 54$\pm$4 | 0.25$\pm$0.03 | 1.0 | 2.29$\pm$0.05 | 0.43 | -1.75 | 0
A19 | 00:01:41.347 | +23:29:33.3 | 4330 | 104$\pm$16 | 0.30$\pm$0.06 | 1.0 | 2.38$\pm$0.08 | 0.43 | -1.75 | 0
A20 | 00:01:41.337 | +23:29:33.1 | 4320 | 97$\pm$8 | 0.39$\pm$0.04 | 1.0 | 2.49$\pm$0.05 | 0.43 | -1.75 | 0
A21 | 00:01:41.370 | +23:29:32.8 | 4462 | 48.3$\pm$12 | 0.10$\pm$0.04 | 0.5 | 1.91$\pm$0.13 | 0.43 | -1.75 | 0
A22 | 00:01:41.378 | +23:29:32.4 | 4505 | 34.2$\pm$3 | 0.16$\pm$0.02 | 1.0 | 2.09$\pm$0.05 | 0.43 | -1.75 | 0
A23 | 00:01:41.188 | +23:29:32.7 | 4505 | 30.0$\pm$2 | 0.10$\pm$0.01 | 1.0 | 1.89$\pm$0.05 | 0.43 | -1.75 | 0
A24 | 00:01:41.200 | +23:29:32.5 | 4520 | 33.5$\pm$3 | 0.14$\pm$0.02 | 1.0 | 2.05$\pm$0.05 | 0.43 | -1.75 | 0
A25 | 00:01:41.281 | +23:29:32.0 | 4342 | 71$\pm$14 | 0.16$\pm$0.05 | 0.7 | 2.11$\pm$0.11 | 0.43 | -1.75 | 0
A26 | 00:01:41.264 | +23:29:31.8 | 4348 | 146$\pm$29 | 0.30$\pm$0.08 | 0.7 | 2.37$\pm$0.11 | 0.43 | -1.75 | 0
A27 | 00:01:41.340 | +23:29:31.2 | 4437 | 73$\pm$15 | 0.07$\pm$0.02 | 0.4 | 1.73$\pm$0.11 | 0.43 | -1.75 | 0
A28 | 00:01:41.219 | +23:29:29.5 | 4469 | 26$\pm$4 | 0.08$\pm$0.02 | 0.5 | 1.78$\pm$0.08 | 0.43 | -1.75 | 0
A29 | 00:01:41.155 | +23:29:28.5 | 4475 | 144$\pm$29 | 0.22$\pm$0.06 | 0.4 | 2.24$\pm$0.11 | 0.43 | -1.75 | 0
A30 | 00:01:41.070 | +23:29:27.2 | 4300 | 68$\pm$14 | 0.14$\pm$0.04 | 0.4 | 2.04$\pm$0.11 | 0.43 | -1.75 | 0
B1 | 00:01:41.476 | +23:29:38.9 | 4500 | 76.9$\pm$6 | 0.53$\pm$0.06 | 1.0 | 2.62$\pm$0.05 | 0.43 | -1.75 | 0
B2 | 00:01:41.466 | +23:29:39.1 | 4477 | 78.7$\pm$6 | 0.47$\pm$0.05 | 1.0 | 2.57$\pm$0.05 | 0.43 | -1.75 | 0
B3 | 00:01:41.450 | +23:29:39.3 | 4466 | 94.4$\pm$8 | 0.41$\pm$0.05 | 1.0 | 2.51$\pm$0.05 | 0.43 | -1.75 | 0
B4 | 00:01:41.455 | +23:29:38.9 | 4481 | 73.0$\pm$6 | 0.38$\pm$0.04 | 1.0 | 2.47$\pm$0.05 | 0.43 | -1.75 | 0
B5 | 00:01:41.421 | +23:29:38.6 | 4380 | 115$\pm$17 | 0.31$\pm$0.06 | 1.0 | 2.39$\pm$0.08 | 0.43 | -1.75 | 0
B6 | 00:01:41.405 | +23:29:38.8 | 4369 | 72$\pm$6 | 0.25$\pm$0.03 | 1.0 | 2.29$\pm$0.05 | 0.43 | -1.75 | 0
B7 | 00:01:41.387 | +23:29:39.5 | 4454 | 78$\pm$6 | 0.49$\pm$0.06 | 1.0 | 2.59$\pm$0.05 | 0.43 | -1.75 | 0
B8 | 00:01:41.373 | +23:29:39.3 | 4449 | 83$\pm$7 | 0.68$\pm$0.08 | 1.0 | 2.73$\pm$0.05 | 0.43 | -1.75 | 0
B9 | 00:01:41.369 | +23:29:39.6 | 4441 | 79$\pm$6 | 0.51$\pm$0.06 | 1.0 | 2.60$\pm$0.05 | 0.43 | -1.75 | 0
B10 | 00:01:41.361 | +23:29:39.9 | 4427 | 61$\pm$5 | 0.26$\pm$0.03 | 1.0 | 2.31$\pm$0.05 | 0.43 | -1.75 | 0
B11 | 00:01:41.310 | +23:29:39.6 | 4428 | 34$\pm$3 | 0.13$\pm$0.01 | 1.0 | 2.01$\pm$0.05 | 0.43 | -1.75 | 0
B12 | 00:01:41.294 | +23:29:39.8 | 4436 | 33$\pm$3 | 0.14$\pm$0.02 | 1.0 | 2.06$\pm$0.05 | 0.43 | -1.75 | 0
B13 | 00:01:41.275 | +23:29:39.7 | 4435 | 31$\pm$2 | 0.14$\pm$0.02 | 1.0 | 2.04$\pm$0.05 | 0.43 | -1.75 | 0
B14 | 00:01:41.304 | +23:29:38.8 | 4354 | 56$\pm$4 | 0.30$\pm$0.03 | 1.0 | 2.38$\pm$0.05 | 0.43 | -1.75 | 0
B15 | 00:01:41.283 | +23:29:38.7 | 4301 | 31$\pm$2 | 0.14$\pm$0.02 | 1.0 | 2.04$\pm$0.05 | 0.43 | -1.75 | 0
B16 | 00:01:41.296 | +23:29:37.7 | 4468 | 116$\pm$23 | 0.30$\pm$0.08 | 0.5 | 2.37$\pm$0.11 | 0.43 | -1.75 | 0
B17 | 00:01:41.279 | +23:29:37.5 | 4462 | 80$\pm$6 | 0.38$\pm$0.04 | 1.0 | 2.48$\pm$0.05 | 0.43 | -1.75 | 0
B18 | 00:01:41.277 | +23:29:37.8 | 4454 | 65$\pm$5 | 0.39$\pm$0.04 | 1.0 | 2.49$\pm$0.05 | 3.07$\pm$0.36 | -0.89 | 1
B19 | 00:01:41.260 | +23:29:37.6 | 4433 | 76$\pm$6 | 0.45$\pm$0.05 | 1.0 | 2.55$\pm$0.05 | 3.58$\pm$0.38 | -0.83 | 1
B20 | 00:01:41.257 | +23:29:38.0 | 4454 | 62$\pm$5 | 0.35$\pm$0.04 | 1.0 | 2.44$\pm$0.05 | 3.43$\pm$0.37 | -0.85 | 1
B21 | 00:01:41.244 | +23:29:38.2 | 4472 | 49$\pm$4 | 0.33$\pm$0.04 | 1.0 | 2.42$\pm$0.05 | 0.43 | -1.75 | 0
B22 | 00:01:41.237 | +23:29:38.5 | 4469 | 47$\pm$4 | 0.33$\pm$0.04 | 1.0 | 2.42$\pm$0.05 | 0.43 | -1.75 | 0
B23 | 00:01:41.236 | +23:29:38.8 | 4454 | 46$\pm$4 | 0.22$\pm$0.03 | 1.0 | 2.25$\pm$0.05 | 0.43 | -1.75 | 0
B24 | 00:01:41.185 | +23:29:38.4 | 4456 | 51$\pm$4 | 0.29$\pm$0.03 | 1.0 | 2.36$\pm$0.05 | 0.43 | -1.75 | 0
B25 | 00:01:41.193 | +23:29:38.7 | 4464 | 67$\pm$5 | 0.24$\pm$0.03 | 1.0 | 2.28$\pm$0.05 | 0.43 | -1.75 | 0
B26 | 00:01:41.174 | +23:29:38.7 | 4466 | 61$\pm$5 | 0.41$\pm$0.05 | 1.0 | 2.51$\pm$0.05 | 0.43 | -1.75 | 0
B27 | 00:01:41.192 | +23:29:39.0 | 4476 | 78$\pm$6 | 0.30$\pm$0.03 | 1.0 | 2.38$\pm$0.05 | 0.43 | -1.75 | 0
B28 | 00:01:41.172 | +23:29:39.0 | 4490 | 65$\pm$5 | 0.36$\pm$0.04 | 1.0 | 2.45$\pm$0.05 | 0.43 | -1.75 | 0
B29 | 00:01:41.158 | +23:29:38.8 | 4496 | 92$\pm$7 | 0.52$\pm$0.06 | 1.0 | 2.61$\pm$0.05 | 0.43 | -1.75 | 0
B30 | 00:01:41.141 | +23:29:38.6 | 4507 | 78$\pm$6 | 0.41$\pm$0.05 | 1.0 | 2.51$\pm$0.05 | 0.43 | -1.75 | 0
B31 | 00:01:41.201 | +23:29:39.3 | 4497 | 68$\pm$5 | 0.36$\pm$0.04 | 1.0 | 2.45$\pm$0.05 | 6.09$\pm$0.47 | -0.60 | 1
B32 | 00:01:41.189 | +23:29:39.5 | 4510 | 80$\pm$12 | 0.25$\pm$0.05 | 1.0 | 2.30$\pm$0.08 | 5.63$\pm$0.46 | -0.63 | 1
B33 | 00:01:41.177 | +23:29:39.3 | 4499 | 56$\pm$4 | 0.25$\pm$0.03 | 1.0 | 2.30$\pm$0.05 | 4.45$\pm$0.41 | -0.73 | 1
B34 | 00:01:41.154 | +23:29:39.1 | 4500 | 73$\pm$6 | 0.36$\pm$0.04 | 1.0 | 2.46$\pm$0.05 | 0.43 | -1.75 | 0
B35 | 00:01:41.139 | +23:29:39.0 | 4480 | 60$\pm$9 | 0.19$\pm$0.04 | 0.5 | 2.18$\pm$0.08 | 0.43 | -1.75 | 0
B36 | 00:01:41.117 | +23:29:38.8 | 4482 | 47$\pm$4 | 0.25$\pm$0.03 | 1.0 | 2.30$\pm$0.05 | 0.43 | -1.75 | 0
C1 | 00:01:41.277 | +23:29:41.1 | 4514 | 42$\pm$6 | 0.13$\pm$0.03 | 1.0 | 2.01$\pm$0.08 | 0.43 | -1.75 | 0
C2 | 00:01:41.276 | +23:29:41.5 | 4521 | 35$\pm$3 | 0.15$\pm$0.02 | 1.0 | 2.06$\pm$0.05 | 0.43 | -1.75 | 0
C3 | 00:01:41.274 | +23:29:41.7 | 4526 | 32$\pm$3 | 0.16$\pm$0.02 | 1.0 | 2.11$\pm$0.05 | 0.43 | -1.75 | 0
C4 | 00:01:41.275 | +23:29:42.0 | 4528 | 25$\pm$2 | 0.08$\pm$0.01 | 1.0 | 1.82$\pm$0.05 | 0.43 | -1.75 | 0
C5 | 00:01:41.234 | +23:29:42.5 | 4568 | 61$\pm$5 | 0.20$\pm$0.02 | 1.0 | 2.21$\pm$0.05 | 0.43 | -1.75 | 0
C6 | 00:01:41.228 | +23:29:42.8 | 4586 | 59$\pm$5 | 0.23$\pm$0.03 | 1.0 | 2.26$\pm$0.05 | 0.43 | -1.75 | 0
C7 | 00:01:41.224 | +23:29:43.0 | 4591 | 42$\pm$3 | 0.21$\pm$0.02 | 1.0 | 2.22$\pm$0.05 | 0.43 | -1.75 | 0
C8 | 00:01:41.207 | +23:29:42.9 | 4584 | 57$\pm$5 | 0.30$\pm$0.03 | 1.0 | 2.38$\pm$0.05 | 0.43 | -1.75 | 0
C9 | 00:01:41.207 | +23:29:43.2 | 4576 | 52$\pm$4 | 0.23$\pm$0.03 | 1.0 | 2.26$\pm$0.05 | 0.43 | -1.75 | 0
C10 | 00:01:41.188 | +23:29:43.0 | 4572 | 44$\pm$4 | 0.21$\pm$0.02 | 1.0 | 2.23$\pm$0.05 | 0.43 | -1.75 | 0
C11 | 00:01:41.208 | +23:29:43.5 | 4583 | 37$\pm$3 | 0.13$\pm$0.01 | 1.0 | 2.00$\pm$0.05 | 0.43 | -1.75 | 0
C12 | 00:01:41.187 | +23:29:43.3 | 4573 | 50$\pm$4 | 0.20$\pm$0.02 | 1.0 | 2.20$\pm$0.05 | 0.43 | -1.75 | 0
C13 | 00:01:41.183 | +23:29:43.6 | 4572 | 42$\pm$6 | 0.12$\pm$0.02 | 0.5 | 1.98$\pm$0.08 | 0.43 | -1.75 | 0
C14 | 00:01:41.136 | +23:29:43.8 | 4578 | 29$\pm$4 | 0.08$\pm$0.02 | 1.0 | 1.81$\pm$0.08 | 0.43 | -1.75 | 0
C15 | 00:01:41.121 | +23:29:44.1 | 4574 | 39$\pm$6 | 0.11$\pm$0.02 | 1.0 | 1.92$\pm$0.08 | 0.43 | -1.75 | 0
C15 | 00:01:41.121 | +23:29:44.1 | 4574 | 39$\pm$6 | 0.11$\pm$0.02 | 1.0 | 1.92$\pm$0.08 | 0.43 | -1.75 | 0
C16 | 00:01:41.087 | +23:29:44.3 | — | — | — | 0.0 | — | 0.43 | -1.75 | 0
C17 | 00:01:41.046 | +23:29:44.2 | — | — | — | 0.0 | — | 0.43 | -1.75 | 0
C18 | 00:01:41.043 | +23:29:44.5 | 4587 | 45$\pm$9 | 0.09$\pm$0.03 | 0.5 | 1.87$\pm$0.11 | 0.43 | -1.75 | 0
D1 | 00:01:41.010 | +23:29:36.6 | 4584 | 25$\pm$2 | 0.07$\pm$0.01 | 1.0 | 1.76$\pm$0.05 | 11.10$\pm$0.62 | -0.34 | 1
D2 | 00:01:41.022 | +23:29:36.9 | 4582 | 30$\pm$2 | 0.19$\pm$0.02 | 1.0 | 2.17$\pm$0.05 | 10.94$\pm$0.62 | -0.34 | 1
D3 | 00:01:40.999 | +23:29:36.9 | 4584 | 25$\pm$2 | 0.07$\pm$0.01 | 1.0 | 1.75$\pm$0.05 | 12.17$\pm$0.65 | -0.30 | 1
D4 | 00:01:41.016 | +23:29:37.3 | 4569 | 48$\pm$4 | 0.18$\pm$0.02 | 1.0 | 2.15$\pm$0.05 | 14.73$\pm$0.71 | -0.21 | 1
D5 | 00:01:40.997 | +23:29:37.4 | 4553 | 25$\pm$2 | 0.13$\pm$0.01 | 1.0 | 2.01$\pm$0.05 | 16.37$\pm$0.74 | -0.17 | 1
D6 | 00:01:40.976 | +23:29:38.0 | 4542 | 73$\pm$6 | 0.36$\pm$0.04 | 1.0 | 2.25$\pm$0.05 | 11.76$\pm$0.64 | -0.31 | 1
D7 | 00:01:40.958 | +23:29:36.3 | 4585 | 38$\pm$3 | 0.26$\pm$0.03 | 1.0 | 2.31$\pm$0.05 | 7.72$\pm$0.53 | -0.49 | 1
D8 | 00:01:40.958 | +23:29:36.7 | 4588 | 39$\pm$3 | 0.26$\pm$0.03 | 1.0 | 2.31$\pm$0.05 | 15.96$\pm$0.73 | -0.18 | 1
D9 | 00:01:40.942 | +23:29:36.6 | 4590 | 43$\pm$3 | 0.43$\pm$0.05 | 1.0 | 2.53$\pm$0.05 | 21.99$\pm$0.86 | -0.04 | 1
D10 | 00:01:40.941 | +23:29:37.0 | 4586 | 51$\pm$4 | 0.37$\pm$0.04 | 1.0 | 2.47$\pm$0.05 | 42.96$\pm$1.18 | 0.25 | 1
D11 | 00:01:40.959 | +23:29:37.3 | 4548 | 25$\pm$2 | 0.18$\pm$0.02 | 1.0 | 2.15$\pm$0.05 | 30.69$\pm$1.01 | 0.11 | 1
D12 | 00:01:40.951 | +23:29:37.5 | 4546 | 27$\pm$2 | 0.22$\pm$0.02 | 1.0 | 2.24$\pm$0.05 | 26.08$\pm$0.93 | 0.03 | 1
D13 | 00:01:40.935 | +23:29:37.5 | 4539 | 36$\pm$3 | 0.28$\pm$0.03 | 1.0 | 2.35$\pm$0.05 | 29.66$\pm$0.99 | 0.09 | 1
D14 | 00:01:40.919 | +23:29:38.1 | 4546 | 38$\pm$3 | 0.22$\pm$0.02 | 1.0 | 2.24$\pm$0.05 | 36.31$\pm$1.09 | 0.18 | 1
D15 | 00:01:40.902 | +23:29:38.0 | 4544 | 35$\pm$3 | 0.27$\pm$0.03 | 1.0 | 2.33$\pm$0.05 | 23.53$\pm$0.88 | -0.01 | 1
D16 | 00:01:40.888 | +23:29:38.1 | 4541 | 47$\pm$4 | 0.29$\pm$0.03 | 1.0 | 2.37$\pm$0.05 | 14.83$\pm$0.71 | -0.21 | 1
D17 | 00:01:40.856 | +23:29:37.6 | 4556 | 36$\pm$3 | 0.14$\pm$0.02 | 1.0 | 2.05$\pm$0.05 | 6.14$\pm$0.48 | -0.59 | 1
D18 | 00:01:40.818 | +23:29:37.7 | 4507 | 90$\pm$7 | 0.36$\pm$0.04 | 1.0 | 2.45$\pm$0.05 | 2.15$\pm$0.31 | -1.05 | 1
D19 | 00:01:40.797 | +23:29:37.5 | 4495 | 75$\pm$15 | 0.19$\pm$0.05 | 1.0 | 2.18$\pm$0.11 | 1.53$\pm$0.28 | -1.20 | 1
D20 | 00:01:40.850 | +23:29:38.0 | 4508 | 71$\pm$6 | 0.25$\pm$0.03 | 1.0 | 2.30$\pm$0.05 | 3.84$\pm$0.39 | -0.80 | 1
D21 | 00:01:40.830 | +23:29:38.0 | 4504 | 62$\pm$5 | 0.32$\pm$0.04 | 1.0 | 2.41$\pm$0.05 | 3.53$\pm$0.38 | -0.83 | 1
D22 | 00:01:40.812 | +23:29:38.0 | 4515 | 70$\pm$6 | 0.31$\pm$0.03 | 1.0 | 2.38$\pm$0.05 | 2.86$\pm$0.35 | -0.92 | 1
D23 | 00:01:40.845 | +23:29:38.3 | 4498 | 59$\pm$5 | 0.23$\pm$0.03 | 1.0 | 2.27$\pm$0.05 | 3.27$\pm$0.37 | -0.87 | 1
D24 | 00:01:40.824 | +23:29:38.3 | 4494 | 50$\pm$4 | 0.22$\pm$0.02 | 1.0 | 2.24$\pm$0.05 | 2.66$\pm$0.34 | -0.96 | 1
D25 | 00:01:40.836 | +23:29:38.8 | 4484 | 46$\pm$4 | 0.19$\pm$0.02 | 1.0 | 2.17$\pm$0.05 | 2.56$\pm$0.33 | -0.97 | 1
D26 | 00:01:40.818 | +23:29:38.6 | 4456 | 91$\pm$18 | 0.24$\pm$0.07 | 1.0 | 2.28$\pm$0.11 | 2.30$\pm$0.32 | -1.02 | 1
D27 | 00:01:40.799 | +23:29:38.5 | 4472 | 59$\pm$5 | 0.23$\pm$0.03 | 1.0 | 2.26$\pm$0.05 | 2.40$\pm$0.33 | -1.00 | 1
D28 | 00:01:40.793 | +23:29:37.0 | 4446 | 25$\pm$2 | 0.09$\pm$0.01 | 1.0 | 1.84$\pm$0.05 | 0.43 | -1.75 | 0
D29 | 00:01:40.771 | +23:29:37.3 | 4489 | 46$\pm$4 | 0.21$\pm$0.02 | 1.0 | 2.22$\pm$0.05 | 0.43 | -1.75 | 0
D30 | 00:01:40.759 | +23:29:37.2 | 4499 | 67$\pm$5 | 0.26$\pm$0.03 | 1.0 | 2.31$\pm$0.05 | 0.43 | -1.75 | 0
D31 | 00:01:40.741 | +23:29:37.2 | 4480 | 69$\pm$6 | 0.26$\pm$0.03 | 1.0 | 2.32$\pm$0.05 | 0.43 | -1.75 | 0
D32 | 00:01:40.723 | +23:29:37.3 | 4467 | 56$\pm$4 | 0.25$\pm$0.03 | 1.0 | 2.30$\pm$0.05 | 0.43 | -1.75 | 0
D33 | 00:01:40.723 | +23:29:36.9 | 4467 | 63$\pm$5 | 0.31$\pm$0.03 | 1.0 | 2.39$\pm$0.05 | 0.43 | -1.75 | 0
D34 | 00:01:40.706 | +23:29:36.6 | 4465 | 72$\pm$11 | 0.23$\pm$0.05 | 1.0 | 2.26$\pm$0.08 | 0.43 | -1.75 | 0
D35 | 00:01:40.746 | +23:29:37.7 | 4476 | 78$\pm$6 | 0.30$\pm$0.03 | 1.0 | 2.37$\pm$0.05 | 0.43 | -1.75 | 0
D36 | 00:01:40.728 | +23:29:37.7 | 4455 | 70$\pm$6 | 0.35$\pm$0.04 | 1.0 | 2.44$\pm$0.05 | 0.43 | -1.75 | 0
D37 | 00:01:40.732 | +23:29:37.9 | 4490 | 104$\pm$8 | 0.42$\pm$0.05 | 1.0 | 2.52$\pm$0.05 | 0.43 | -1.75 | 0
D38 | 00:01:40.715 | +23:29:37.8 | 4467 | 78$\pm$6 | 0.29$\pm$0.03 | 1.0 | 2.37$\pm$0.05 | 0.43 | -1.75 | 0
D39 | 00:01:40.730 | +23:29:38.3 | 4512 | 69$\pm$6 | 0.28$\pm$0.03 | 1.0 | 2.34$\pm$0.05 | 7.31$\pm$0.51 | -0.52 | 1
D40 | 00:01:40.659 | +23:29:36.4 | 4429 | 53$\pm$8 | 0.17$\pm$0.03 | 1.0 | 2.13$\pm$0.08 | 0.43 | -1.75 | 0
D41 | 00:01:40.674 | +23:29:35.6 | 4347 | 52$\pm$8 | 0.16$\pm$0.03 | 1.0 | 2.10$\pm$0.08 | 0.43 | -1.75 | 0
E1 | 00:01:40.875 | +23:29:37.1 | 4509 | 27$\pm$4 | 0.09$\pm$0.02 | 1.0 | 1.83$\pm$0.08 | 14.86 | -0.21 | 1
E2 | 00:01:40.870 | +23:29:36.6 | 4463 | 57$\pm$5 | 0.21$\pm$0.02 | 1.0 | 2.21$\pm$0.05 | 0.43 | -1.75 | 0
E3 | 00:01:40.851 | +23:29:36.5 | 4465 | 45$\pm$4 | 0.19$\pm$0.02 | 1.0 | 2.18$\pm$0.05 | 0.43 | -1.75 | 0
E4 | 00:01:40.852 | +23:29:36.1 | 4445 | 57$\pm$5 | 0.21$\pm$0.02 | 1.0 | 2.23$\pm$0.05 | 0.43 | -1.75 | 0
E5 | 00:01:40.859 | +23:29:35.8 | 4402 | 75$\pm$11 | 0.23$\pm$0.05 | 1.0 | 2.26$\pm$0.08 | 0.43 | -1.75 | 0