# ALMA Observation of NGC 5135: The Circumnuclear
${\mathrm{CO}}\,(6\text{--}5)$ and Dust Continuum Emission at 45 Parsec
Resolution\star\star$\star$ The National Radio Astronomy Observatory is a
facility of the National Science Foundation operated under cooperative
agreement by Associated Universities, Inc.

Tianwen Cao 11affiliation: Chinese Academy of Sciences South America Center
for Astronomy, National Astronomical Observatories, Chinese Academy of
Sciences, Beijing 100101, China<EMAIL_ADDRESS>22affiliation: Key Laboratory
of Optical Astronomy,National Astronomical Observatories, Chinese Academy of
Sciences, Beijing 100101, China 33affiliation: School of Astronomy and Space
Science, University of Chinese Academy of Sciences, Beijing 100049, China
44affiliation: Instituto de Astrofisica, Pontificia Universidad Católica de
Chile,Av.Vicuña Mackenna4860, 7820436 Macul,Santiago, Chile , Nanyao Lu
11affiliation: Chinese Academy of Sciences South America Center for Astronomy,
National Astronomical Observatories, Chinese Academy of Sciences, Beijing
100101, China<EMAIL_ADDRESS>, C. Kevin Xu 11affiliation: Chinese Academy of
Sciences South America Center for Astronomy, National Astronomical
Observatories, Chinese Academy of Sciences, Beijing 100101, China;
<EMAIL_ADDRESS>, Yinghe Zhao 55affiliation: Yunnan Observatories, Chinese
Academy of Sciences, Kunming 650011, China 66affiliation: Key Laboratory for
the Structure and Evolution of Celestial Objects, Chinese Academy of
Sciences,Kunming 650011, China 77affiliation: Center for Astronomical Mega-
Science, CAS, 20A Datun Road, Chaoyang District, Beijing 100012, China , Venu
Madhav Kalari 88affiliation: Departamento de Astronomía, Universidad de Chile,
Casilla 36-D, Correo Central,Santiago,Chile , Yu Gao 99affiliation: Purple
Mountain Observatory/Key Lab of Radio Astronomy, 8 YuanHua Road, 210034
Nanjing, China , Vassilis Charmandaris 1010affiliation: Department of Physics,
University of Crete, GR-71003 Her- aklion, Greece 1111affiliation: IAASARS,
National Observatory of Athens, GR-15236, Pen-teli, Greece , Tanio Diaz Santos
1212affiliation: Núcleo de Astronomía de laFacultad deIngeniería, Universidad
Diego Portales,Av. Ejercito Libertador 441,Santiago,Chile , Paul van der Werf
1313affiliation: Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA
Leiden, The Netherlands , Chen Cao 1414affiliation: School of Space Science
and Physics, Shandong University, Weihai, Weihai, Shandong 264209, China
1515affiliation: Shandong Provincial Key Laboratory of Optical Astronomy &
Solar-TerrestrialEnvironment, Weihai, Shandong 264209,China , Hong Wu
22affiliation: Key Laboratory of Optical Astronomy,National Astronomical
Observatories, Chinese Academy of Sciences, Beijing 100101, China
33affiliation: School of Astronomy and Space Science, University of Chinese
Academy of Sciences, Beijing 100049, China , Hanae Inami 1616affiliation:
Univ. Lyon, Univ. Lyon1, ENS de Lyon, CNRS, Centre de Recherche Astrophysique
de Lyon (CRAL) UMR 5574, 69230 Saint-Genis-Laval, France , Aaron Evans
1717affiliation: Department of Astronomy, University of Virginia,
Charlottesville, VA 22903-2325, USA 1818affiliation: National Radio Astronomy
Observatory, 520 Edgemont Road, Charlottesville, VA 22903-2475, USA

###### Abstract

We present high-resolution (0.17″$\times$ 0.14″) Atacama Large
Millimeter/submillimeter Array (ALMA) observations of the
${\mathrm{CO}}\,(6\text{--}5)$ line, and 435$\,\mu{\rm m}$ dust continuum
emission within a $\sim$9″$\times$ 9″ area centered on the nucleus of the
galaxy NGC 5135. NGC 5135 is a well-studied luminous infrared galaxy that also
harbors a Compton-thick active galactic nucleus (AGN). At the achieved
resolution of 48 $\times$ 40 pc, the ${\mathrm{CO}}\,(6\text{--}5)$ and dust
emissions are resolved into gas “clumps” along the symmetrical dust lanes
associated with the inner stellar bar. The clumps have radii between
$\sim$45-180 pc and ${\mathrm{CO}}\,(6\text{--}5)$ line widths of $\sim$60-88
km s-1. The ${\mathrm{CO}}\,(6\text{--}5)$ to dust continuum flux ratios vary
among the clumps and show an increasing trend with the [Fe ii]/Br-$\gamma$
ratios, which we interpret as evidence for supernova-driven shocked gas
providing a significant contribution to the ${\mathrm{CO}}\,(6\text{--}5)$
emission. The central AGN is undetected in continuum, nor in CO (6-5) if its
line velocity width is no less than $\sim$ 40  km s-1. We estimate that the
AGN contributes at most 1% of the integrated ${\mathrm{CO}}\,(6\text{--}5)$
flux of 512 $\pm$ 24 Jy km s-1 within the ALMA field of view, which in turn
accounts for $\sim$32% of the ${\mathrm{CO}}\,(6\text{--}5)$ flux of the whole
galaxy.

###### Subject headings:

galaxies: active – galaxies: starbusrt – galaxies: general – galaxies: nuclei
– galaxies: ISM

††slugcomment: Accepted for publication in APJ;

##  1\. Introduction

Luminous infrared galaxies (LIRGs; $L_{\rm IR}^{[8-1000\,\mu\rm m]}\gtrsim
10^{11}\,L_{\odot}$), whose space density exceeds that of optically selected
starburst and Active Galactic Nucleus (AGN) host galaxies at comparable
bolometric luminosities (Soifer et al., 1987), consist of isolated galaxies,
galaxy pairs, interacting galaxy systems and advanced mergers (Sanders &
Mirafbel, 1996; Wang et al., 2006). LIRGs in the later stages of evolution
tend to contain a rich amount of molecular gas in the galaxy nuclear region
(Sanders et al., 1986), and have a higher fraction of AGNs compared with less
luminous galaxies (Sanders & Mirafbel, 1996). Detailed investigations of the
physical properties, AGN-starburst connection, and gas inflow/outflow in
representative LIRGs in the local Universe are critical to our understanding
of galaxy evolution because LIRGs are the dominant contributors to the cosmic
star formation (SF) at $z\gtrsim 1$ (Le Floc’h et al. 2005; Caputi et al.
2007; Magnelli et al. 2009, 2011; Gruppioni et al. 2013)

The CO emission lines from low-$J$ transitions, such as CO (1-0) at 2.6 mm and
CO (2-1) at 1.3 mm, have been widely used to trace the molecular gas content
in LIRGs (Solomon & Sage, 1988; Sanders et al., 1991; Solomon et al., 1997;
Bryant & Scoville, 1999; Gao & Solomon, 1999; Evans et al., 2002). However,
based on the data taken with the SPIRE Fourier Transform Spectrometer (FTS;
Griffin et al. 2010) onboard the Herschel Space Observatory (Herschel;
Pilbratt et al. 2010) on a flux-limited sample of 123 LIRGs from the Great
Observatories All-Sky LIRGs Survey (GOALS; Armus et al. 2009), Lu et al.
(2014, 2017) showed that the mid-$J$ CO emission (i.e., 4 $<J<$ 10) from warm
and dense molecular gas correlates linearly with the star formation rate (SFR)
on galactic scale for LIRGs over a wide range of $L_{\rm IR}$ and far-infrared
(FIR) color. Therefore, the heating mechanism for the warm dense gas that
gives rise to the mid-$J$ CO line emission should ultimately derive the energy
from the same SF process that powers the dust emission. There is not yet a
firm consensus on this heating mechanism. Apparently different heating
mechanisms are favored from analyses of the CO emission line spectra of
individual galaxies, including far-UV photon heating (e.g., Rigopoulou et al.
2013), heating by cosmic rays enhanced by supernovae (SNe; e.g., Bradford et
al. 2003), and heating by shocks that may or may not be powered by SNe (e.g.,
Kamenetzky et al. 2012; Rangwala et al. 2011; Nikola et al. 2011; Meijerink et
al. 2013; Pellegrini et al. 2013; Rosenberg et al. 2014). The X-ray photons
from an AGN can heat the surrounding dense gas very effectively (e.g., Spaans
& Meijerink 2008). However, Lu et al. (2017) argued that the CO line emission
associated with any AGN gas heating may peak at $J>10$. As a result, the
mid-$J$ CO line emission is always dominated by SF.

With the Atacama Large Millimeter/submillimeter Array (ALMA; Wootten &
Thompson 2009), it is now possible to obtain high-resolution mid-$J$ CO line
and dust continuum images of the nuclei of nearby LIRGs to investigate whether
the Herschel results above still hold true at physical scales down to the
typical size of giant molecular clouds (GMCs; i.e., $\sim$40 pc, Kawamura et
al. 2009), and whether one can rule out some of the gas heating mechanisms
proposed. To this end, we have carried out a number of ALMA Band-9
observations over time, to image simultaneously the CO (6-5) line emission
(the rest frequency $\nu_{\rm rest}$ = 691.473 GHz) and its underlying dust
continuum at 435$\,\mu{\rm m}$ in the nuclear regions of a set of carefully
selected, representative LIRGs from our Herschel FTS sample. The targets
observed include NGC 34 (Xu et al., 2014) and NGC 1614 (Xu et al., 2015), two
advanced mergers with a warm FIR color; NGC 7130 (Zhao et al., 2016) and NGC
5135 (in this paper), two well-known Seyfert galaxies with a prominent stellar
bar; IC 5179 (Zhao et al., 2017), an isolated, unbarred galaxy with a compact
nuclear starburst; and CGCG 049-057 with a high-surface density nuclear SF
disk (Cao et al., in preparation). The linear resolutions ($R_{\rm linear}$)
achieved range from 100 pc in the early observation of NGC 34 to 34 pc in the
case of IC 5179. Three additional LIRGs in our FTS sample also have ALMA
${\mathrm{CO}}\,(6\text{--}5)$ images in the literature: Arp 220 ($R_{\rm
linear}\sim 165$ pc; Wilson et al. 2014; Rangwala et al. 2015), IRAS
13120-5453 ($R_{\rm linear}\sim 165$ pc; Sliwa et al. 2017), and NGC 1068
($R_{\rm linear}\sim 4$ pc; García-Burillo et al. 2014, 2016). ALMA CO (6-5)
images also exist for two nearby, but non-LIRG galaxies: NGC 1377 (Aalto et
al., 2017) and Centaurus A (Espada et al., 2017).

At a distance of 59 Mpc (1″ corresponds to 281 pc) and with a fairly face-on
disk, NGC 5135 is a well studied LIRG with $L_{\rm IR}=10^{11.33}\,L_{\odot}$
and a moderately warm FIR color of 0.54 (in terms of the 60-to-100 $\,\mu{\rm
m}$ flux density ratio; Armus et al. 2009). The galaxy not only displays a
powerful circumnuclear starburst over a region of $\sim$1 kpc in diameter
(González Delgado et al., 1998; Bedregal et al., 2009) but also harbors a
highly obscured Seyfert 2 nucleus (Levenson et al., 2004; Phillips et al.,
1983; Turner et al., 1997). It is therefore an ideal target for high-
resolution ALMA observations to separate the circumnuclear SF from the AGN.
The 6 cm radio continuum emission peaks in an area $\sim$3″ south of the
nucleus based on a Very Large Array (VLA) observation by Ulvestad & Wilson
(1989), presumably tracing the supernova remnants (SNRs) from a previous
starburst. The high-resolution Hubble Space Telescope (HST) UV/optical imaging
observations unveiled a large number of young star clusters, between the
nucleus and the radio continuum peak (González Delgado et al., 1998), which
presumably have partially cleared gas. Furthermore, the high-resolution near-
and mid-infrared images (Alonso-Herrero et al., 2006; Díaz-Santos et al.,
2008) show patches of strong on-going SF along, but at the downstream side of
the dust lanes that are likely associated with the stellar bar (e.g. (Mulchaey
& Regan, 1997)). An intermediate resolution ($R\sim$ 3000$-$4000), near-
infrared integral-field spectroscopy (Bedregal et al., 2009) confirmed the
presence of a high-excitation ionization cone centered on the AGN, based on
the [Si vi] 1.96$\,\mu{\rm m}$ line emission, as well as an extended
distribution of shocked gas likely powered by SNe, based on the [Fe ii]
1.46$\,\mu{\rm m}$ line. Fukazawa et al. (2011) and Singh et al. (2012)
obtained broad-band (10$-$50 kev) X-ray spectra of NGC 5135, demonstrating
that the AGN in NGC 5135 is obscured by Compton-thick material. Our ALMA
imaging of NGC 5135 presented here provides for the first time the
distribution and kinematics of the warm and dense molecular gas as well as the
morphology of the 435$\,\mu{\rm m}$ dust emission, at a linear resolution of
less than 50 pc in the circumnuclear region of NGC 5135.

In the remainder of the paper, we describe our ALMA observation and data
reduction in §2 and present our results in §3. In §4, we discuss the physical
implications derived from our data on the circumnuclear SF and the role of the
AGN, compare our ALMA images with existing images at other wavelengths, and
comment on the most likely heating mechanism for the observed
${\mathrm{CO}}\,(6\text{--}5)$ emission. Thereby allowing us to distinguish
between different heating mechanisms, and the role played by SF in giving rise
to the mid-J CO emission. Finally, we summarize our results in §5. Throughout
this paper, we adopt a distance of 59 Mpc for NGC 5135 (Armus et al. 2009).

##  2\. OBSERVATION AND DATA REDUCTION

The ALMA band-9 observation of NGC 5135 was carried out in the time division
mode (with a velocity resolution of $\sim$6.8  km s-1). The four basebands
(i.e., Spectral Windows; SPWs 0-3) were centered on sky frequencies of
681.975, 683.736, 678.243 and 680.183 GHz, respectively, with a bandwidth of
1.875 GHz. The observation was performed with the configuration mode C34-5,
using 39 12-meter antennae with baselines ranging from 21.3 to 885.6 meters.
The total on-target integration time is 21.03 minutes. During the observation,
the phase calibration and amplitude were monitored using J1316-3328.
Additional observing details can been found in Table 1.

The data were reduced using the Common Astronomical Software Application
(CASA) version 4.5 (McMullin et al., 2007). Our primary beam is $\sim$8.8′′
and the Maximum Recoverable Scale
(MRS)111https://science.nrao.edu/science/videos/largest-angular-scale-and-
maximum-recoverable-scale is 3.5″. The ${\mathrm{CO}}\,(6\text{--}5)$ line
data cube was generated using the data in the SPW covering the sky frequency
range of 680.975 to 682.975 GHz. The continuum was estimated by combining the
data from the other three SPWs. The calibrated images were cleaned using the
Briggs weighting (with the parameter “robust” set to 0.5). The resulting
synthesized beams are nearly identical between the continuum and the line
emission and have a full width at half maximum (FWHM) size of $\sim$ 0.17″
$\times$ 0.14″(equivalent to 48 $\times$ 40 pc at the distance to NGC 5135),
with the major-axis position angle (north to east) at 111°. All our analyses
in this paper use the data after the primary beam correction, whereas all of
the figures (except Fig. 6) were produced using data prior to the primary beam
correction.

The final spectral cube has a channel width equivalent to 13.5 km s-1 in
velocity. The channel noise ($\sigma_{\rm ch}$) is on the order of 18 mJy
beam-1. The total CO (6-5) flux image, as an integration over the barycentric
velocities from 3,971 to 4,157 km s-1, has an r.m.s. noise of $\sim$1.2 Jy
beam-1 km s-1. The continuum image has a noise of 2.2 mJy beam-1. All these
noise measurements were done on the data without the primary beam correction.
The ALMA absolute flux calibration is estimated to be good to $\sim$10%. The
astrometric accuracy is better than 0.01″.

##  3\. Results

###  3.1. CO (6-5) Line Emission

The four panels in Fig. 1 show respectively the images of the total CO (6-5)
emission integrated between the observed velocities of 3971 and 4157  km s-1,
the 435 $\,\mu{\rm m}$ continuum, the velocity field (i.e., moment 1), and
velocity dispersion map (moment 2). Each image roughly covers the ALMA primary
beam. The contours overlaid in Fig. 1a and 1b refer to the same total CO (6-5)
emission and start at S/N $=3$.

The CO (6-5) emission detected at S/N $>3$ appears clumpy and is confined to a
few discrete regions along two spiral arm-like features corresponding
spatially to the dust lanes seen in the UV and optical (e.g., Muñoz Marín et
al. 2007). We mark four separate regions, namely (a, b, c, d) which are
compact and concentrated in the intergrated CO (6-5) map. The region d appears
more diffuse compared with the other three. The total CO (6-5) flux from the
combined four regions in Fig. 1a is 512 $\pm$ 24 Jy km s-1. With a much larger
beam of $\sim$31″, the Herschel/FTS observation gives a
${\mathrm{CO}}\,(6\text{--}5)$ flux of 1,617 Jy km s-1(Lu et al., 2017).
Therefore the clump regions in Fig. 1a together account for $\sim$32% of the
total CO (6-5) flux of the galaxy. The ′′missing′′ line flux could be due to a
combination of the line emission outside the ALMA field of view, or possible
faint emission at peak surface brightness below our 3-$\sigma$ (i.e., 3.6 Jy
beam-1 km s-1) detection limit, or resolved out on larger scales. We analyze
individual clumps in more detail in §4.2.

We set the threshold at 4-$\sigma_{\rm ch}$, where $\sigma_{\rm ch}$ is the
r.m.s noise per frequency channel, to obtain the moment 1 and 2 maps. To
reveal the kinematics better, the moment 1 and 2 maps, shown respectively in
Fig. 1c and 1d, were based on the uv-taper image (the details about the uv-
taper image are presented in the last paragraph of this subsection). The
velocity scale in Fig. 1c was calculated using the formula $\nu_{\rm
obs}=\nu_{\rm rest}\,(1-V/c)$, where $\nu_{\rm obs}$ is the observed CO (6-5)
line frequency, $c$ the speed of light, and $V$ the velocity to calculate. The
line velocity ranges from 3992 to 4140  km s-1. Fig. 1d shows that the line-
of-sight velocity dispersion ranges from 10 to 40  km s-1, using only those
pixels with S/N $>$ 4-$\sigma_{\rm ch}$. The overall kinematic pattern can
also be seen in the channel maps displayed in Fig. 2, where the contours from
an individual channel of width 13.5  km s-1are overlaid on the grayscale image
of the total CO (6-5) flux map shown in Fig. 1a. While the regions a, c, d are
mainly confined within a velocity range of 3992 to 4073  km s-1, the region b
has a range between 4046 and 4127  km s-1. This suggests that the observed
velocity pattern is not dominated by a simple rotation within the galaxy disk.
In the channel maps, some CO (6-5) clumps break down into smaller clumps (or
clouds) in some velocity channels, e.g., the region a in the channel centered
at $V=4073$ km s-1. These clouds have sizes $<50$ pc.

In Fig. 3a and 3b, we reproduced the same images as in Fig. 1a and 1b,
respectively, but with a larger effective beam of 0.4′′ $\times$ 0.4′′
(equivalent to 112 $\times$ 112 pc) by applying an uv-taper (with the
parameter “outertaper” = 0.4′′) to our uv data before imaging. The peak signal
of the region d is higher than 4-$\sigma$ ($\sigma$ = 8 Jy beam-1 km s-1). The
total flux of the four regions combined, as defined in Fig. 1, is 841 $\pm$ 45
Jy km s-1. This flux equals 1.6 times the flux from the original ALMA image,
and is $\sim$52$\%$ of the total flux measured by Herschel, confirming that
there exists some more diffuse or lower surface brightness CO (6-5) emission
within the region of Fig. 1a.

###  3.2. Dust Continuum Emission

As shown in Fig. 3b, the continuum at 435$\,\mu{\rm m}$ generally coincides
with the CO (6-5) line emission in regions a and c at scales of 0.4″
(equivalent to 112 pc), which corresponds to an angular size of 114 pc at the
distance of NGC 5135. This is consistent with the findings in the other LIRGs
we imaged in CO (6-5), i.e., NGC 34, NGC 1614 and NGC 7130 and IC 5179 (Xu et
al., 2014, 2015; Zhao et al., 2016, 2017), i.e., at scales $\gtrsim$ 100 pc,
there is a good spatial correspondence between the CO (6-5) line and its
underlying continuum emissions.

However, at scales significantly smaller than 100 pc, there are apparent
offsets between the local peaks of the line and continuum emissions in Fig.
1b. This small-scale offset between the line and continuum emissions is also
seen in the LIRGs of moderately high nuclear gas surface densities, e.g., IC
5179 (at linear resolution $R_{\rm linear}\approx 34$ pc; Zhao et al. 2017)
and NGC 7130 ($R_{\rm linear}\approx 70$ pc $\times$ 40 pc; Zhao et al. 2016).
Furthermore, in both Fig. 1b and 3b, the dust continuum is unusually weak
relative to the line emission in the regions b and d. As argued in Zhao et al.
(2016), these differences between the line and dust continuum emissions at
small scales can only be understood if the gas and dust are heated by
different mechanisms. We discuss this in more detail in §4.3.

By combining the four regions in Fig. 1b, we derived a total flux of 181 $\pm$
25 mJy for the 435$\,\mu{\rm m}$ continuum emission. This flux would be 1.6
times higher if we had derived it from the same regions in Fig. 3b.

##  4\. ANALYSIS AND DISCUSSION

###  4.1. The Central AGN

The AGN position can be constrained by the peak of the [Si vi] line emission
at 1.96 $\,\mu{\rm m}$ in a ground-based observation (Bedregal et al., 2009)
and by the peak of the hard X-ray (4-8 keV) emission detected with Chandra
(Levenson et al., 2004). The estimated astrometric uncertainty associated with
either of these images is on the order of 0.5″. The VLA 6 cm radio continuum
image of Ulvestad & Wilson (1989) has a modest resolution of 0.91″ $\times$
0.60″ and an astrometric accuracy of 0.3″. We overlaid our CO (6-5) contours
from Fig. 1a on this radio continuum image in Fig. 4a. As already stated in
Sec. 1, the main peak of the 6 cm emission is $\sim$3″ south of the galaxy
nucleus, spatially coincident with the peak of the broad (FWHM = $\sim$513 km
s-1) [FeII] emission (Bedregal et al., 2009). However, there is a minor radio
emission peak (10$\sigma$) near the anticipated AGN position. We take the
position of this radio peak (R.A. = 13h25m44s.02, Dec. = $-$29∘50′00′′.4;
J2000) as the AGN location (i.e., marked by the white cross), with a
positional uncertainty of 0.3-0.5″. In Fig. 4b, we overlaid the same CO (6-5)
contours on the hard X-ray emission.

As shown in Fig. 1a, the ${\mathrm{CO}}\,(6\text{--}5)$ emission is undetected
at 4-$\sigma$ level at the AGN position (i.e., integrated over the velocity
range of 186  km s-1). This would hold true even if we had lowered the
detection threshold to 3-$\sigma$. We assume the AGN-related
${\mathrm{CO}}\,(6\text{--}5)$ emission is confined to an area smaller than
our ALMA beam size ($\sim$48 $\times$ 40 pc), then the 3-$\sigma$ flux upper
limit is equal to 3 $\times$ (1.2 Jy  km s-1) = 3.6 Jy  km s-1($\sigma$ = 1.2
Jy beam${}^{-1}\,$ km s-1).

However, an apparent narrow emission feature at the AGN location is seen at
3-4 $\sigma$ significance over only two velocity channels (i.e., V = 4019.0
and 4032.5  km s-1). The image summed over these two velocity channels is
presented in Fig. 5a, which shows a peak surface brightness of 33 mJy beam-1
(at 5-$\sigma$ significance). The spectrum in Fig. 5b is extracted from a
circular aperture of radius = 0.4′′ (= 2.5 times the FWHM of the 3-$\sigma$
surface brightness of the emission in Fig. 5a). Its narrow velocity width of
$\sim$40  km s-1makes it unlikely that this signal is physically associated
with the AGN. Nevertheless, considering that the
${\mathrm{CO}}\,(6\text{--}5)$ emission associated with the gas torus of the
AGN in NGC 1068 is observed to have only a modest velocity width of $\sim$80
km s-1(García-Burillo et al., 2016), we defer to a future observation of
higher angular resolution to firmly conclude the reality of this narrow
${\mathrm{CO}}\,(6\text{--}5)$ emission. Flux-wise, the narrow
${\mathrm{CO}}\,(6\text{--}5)$ emission in Fig. 5b has a flux of 1.7 Jy  km
s-1, which is smaller than the 3-$\sigma$ flux upper limit of 3.6 Jy  km
s-1derived above. Therefore, we conclude that the AGN in NGC 5135 contributes
at most 1% of the ${\mathrm{CO}}\,(6\text{--}5)$ flux observed within the ALMA
field of view. This is consistent with the Herschel finding that the mid-$J$
CO line emission in LIRGs is mainly associated with SF regardless of whether
there is an AGN or not (Lu et al., 2017). The fractional contribution of the
AGN to the bolometric luminosity of NGC 5135 is about (24$\pm$6)% (Díaz-Santos
et al., 2017). The AGN in NGC 5135 is heavily obscured; the surrounding gas
could be heated to a very high temperature by the X-rays associated with the
AGN, resulting in a CO spectral line distribution that peaks $J>10$ (Spaans &
Meijerink, 2008). Such a scenario seems to be the case in the Seyfert galaxy
NGC 1068: While the high-resolution ALMA imaging shows that the vast majority
of the nuclear CO (6-5) emission is associated with the compact circumnuclear
ring of SF at a radius of $\sim$100 pc (García-Burillo et al., 2016), the
total nuclear CO emission line spectrum has a distinct component that peaks at
$J\sim 16$ (Spinoglio et al., 2012). This hot spectral component of the CO
emission is presumably due to the AGN in NGC 1068. The central AGN in NGC 5135
is bright in terms of the 1-0 S(1) 2.12 $\,\mu{\rm m}$ H2 ro-vibrational line
(Bedregal et al., 2009) associated with warm molecular gas. Although this line
could be excited by different physical processes, including UV-fluorescence
(photons); shock fronts (collisions) and X-ray illumination, Bedregal et al.
(2009) argued that the excited near-IR H2 emission is mainly caused by X-ray
illumination in the AGN region of NGC 5135. Such an X-ray dominant scenario is
also favored based on non-detection of the ${\mathrm{CO}}\,(6\text{--}5)$
emission here.

The 435$\,\mu{\rm m}$ dust continuum is also undetected at the AGN position,
with the 3-$\sigma$ flux upper limit equal to 5.4 mJy ($\sigma$ = 1.8 mJy
beam-1; the same method used for deriving the ${\mathrm{CO}}\,(6\text{--}5)$
flux upper limit). We compare this flux upper limit with the expected
435$\,\mu{\rm m}$ continuum flux from an average infrared spectral energy
distribution (SED) appropriate for the AGNs of X-ray luminosities comparable
to that of NGC 5135: The intrinsic 2.0-10 keV X-ray luminosity of NGC 5135 is
$\sim$1.8 $\times$ 1043 erg s-1 (Singh et al., 2012). We therefore used the
infrared AGN SED for $L_{2.0-10\,\,\rm keV}>10^{42.9}$ erg s-1 in Mullaney et
al. (2011) and anchored it at the 12$\,\mu{\rm m}$ luminosity of NGC 5135
estimated from the X-ray and mid-IR correlation given in Asmus et al. (2015).
This derived SED is shown in Fig. 6, along with two continuum flux upper
limits (at 3$\sigma$) at 435 and 1,300$\,\mu{\rm m}$ based on the ALMA
observation. The latter continuum flux upper limit was estimated from an
archival ALMA Band-6 observation (Project 2013.1.00243.S; PI: L. Colina). This
plot suggests that the ALMA data points are consistent with what is expected
from the typical infrared SED for AGNs like NGC 5135.

###  4.2. Properties of Molecular Gas Clumps

Several clumpy features are resolved in the ${\mathrm{CO}}\,(6\text{--}5)$
image shown in Fig. 1a. The resolved clump sizes of $\sim$100 pc are
comparable with or larger than the beam size. For other (U)LIRGS, such clumpy
features are traced more commonly by low-J CO or isotopologues, owing to the
difficulty in observing dense tracers. However, it is more appropriate to
analyse the properties of compact clumpy structures as seen in Fig. 1a using
denser gas tracers rather than the diffuse gas traced by low-J CO observation.
Dense gas tracers, such as ${\mathrm{CS}}\,(2\text{--}1)$ , HCN (1-0) and
${\mathrm{CO}}\,(6\text{--}5)$ usually trace embedded cloud clumps or cores
within a more extended distribution of CO emission (Leroy et al., 2015;
Sakamoto et al., 2011; Rosolowsky & Blitz, 2005), and therefore are very
useful for studying the dense, embedded star-forming structures within a much
larger molecular region.

In Fig. 1a, the ${\mathrm{CO}}\,(6\text{--}5)$ emission peaks are resolved
into separate clumps at S/N = 4, labelled as a1, a2, a3, a4, a5, b, c and d.
For each clump, we list in Table 2 a number of parameters derived from the
image in Fig. 1a for all the clumps except for the clump d. At the resolution
of Fig. 1a, the clump d is detected only at S/N = 3 and appears to be quite
diffuse. We therefore derived its parameters from the image in Fig. 3a, which
has a coarser resolution of 112 $\times$ 112 pc.

The size of a clump is specified by its FWHM major and minor axes plus the
major axis position angle (PA), which are respectively given in Columns (2)
and (3) of Table 2. These were derived from a 2d Gaussian fit to the clump
intensity map. For the blended clumps (a4 and a5), we segmented the cloud into
sub-clouds employing a variant of the CLUMPFIND algorithm (Williams et al.,
1994). We divided the blended clumps by the half distance of two peaks and
measure their parameters. We also calculated the effective clump radius $R$,
following Solomon et al. (1987). After the deconvolution with the appropriate
ALMA beam (using the CASA function “deconvolvefrombeam”), the radii range from
45 to 180 pc for the clumps (see Table 2, Column (4)).

We also extracted the 1d spectrum for each clump within an elliptical aperture
with radii of (major and minor axes (FWHM)) as the oval areas marked in Fig.
1a. The resulting spectra are plotted for all the clumps in Fig. 7. Using the
same elliptical aperture, we derived the integrated CO (6-5) flux from the
clump intensity map and 435$\,\mu{\rm m}$ continuum flux density from
continuum map. The line central velocity, the line velocity width $\Delta
V_{\rm FWHM}$, the integrated CO (6-5) flux and 435$\,\mu{\rm m}$ continuum
flux density given respectively in Columns (5), (6), (8) and (9) in Table 2.
The resulting $\Delta V_{\rm FWHM}$ ranges from 60 to 88  km s-1.

We also estimated the Virial and molecular gas masses for each clump.
Following Larson (1981); Heyer et al. (2009), the Virial mass ($M_{vir}$) is
estimated as

$M_{\rm vir}/M_{\odot}=\frac{5\,(\Delta V_{\rm
FWHM}/{\rm\mbox{\,km~{}s${}^{-1}$}})^{2}\,R/{\rm pc}}{G},$ (1)

where $R$ is simply the effective clump radius and $\Delta V_{\rm FWHM}$
refers to the value of after deconvolution with channel width 13.5  km s-1and
a possible contribution of about 10 km s-1 from the disk rotation is further
removed. (This correction amount was set to the mean velocity change over the
size of one ALMA beam by examining the P-V plots of all the clumps. In the
Following, the $\Delta V_{\rm FWHM}$ which we have used in Fig. 8 and 9 are
those corrected ones.) The derived Virial masses, shown in Column (10) of
Table 2, rage from $\sim$7 to 60 $\times 10^{7}\,M_{\odot}$.

We estimated the molecular gas mass of a clump, $M_{\rm mol}$, using the
850$\,\mu{\rm m}$ continuum flux density-based formula in Scoville et al.
(2016) by converting the observed 435$\,\mu{\rm m}$ flux density to that at
850$\,\mu{\rm m}$ assuming a dust temperature of the 25 K:

$\displaystyle L_{\nu_{850\mbox{$\,\mu{\rm m}$}}}$
$\displaystyle=1.19\\!\times\\!10^{27}\\!\times\\!S_{\nu}/{\rm
Jy}\\!\times\\!\frac{(\nu_{850\mbox{$\,\mu{\rm
m}$}})^{3.8}}{\nu_{obs}(1+z)}\\!\times\\!\frac{(d_{L}/{\rm Mpc})^{2}}{1+z}$
(2) $\displaystyle\\!\times\\!\frac{\Gamma_{\rm
RJ}(25,\nu_{850\mbox{$\,\mu{\rm m}$}},0)}{\Gamma_{\rm RJ}(25,\nu_{\rm
obs},z)}/(\rm erg\,s^{-1}Hz^{-1})$

and

$\alpha_{\nu}\\!\\!=\\!\\!L_{\nu_{850\\!\\!\mbox{$\,\mu{\rm m}$}}}/M_{\rm
mol}\\!\\!=\\!\\!6.7\pm 1.7\\!\times\\!10^{19}/(\rm
erg\,s^{-1}Hz^{-1}M_{\odot}^{-1})\\\ $

Where $\Gamma_{\rm RJ}(T_{\rm d},\nu,z)=\frac{h\nu/kT_{\rm d}}{e^{h\nu/kT_{\rm
d}}-1}$ and the $S_{\nu}$ is the dust emission flux density. As the
435$\,\mu{\rm m}$ flux density shown in Table 2 is measured within FWHM
(diameter), it should only account for about 58% of the total flux. Thus, we
multiply the flux density by 1.731 to estimate the total flux. For the gas
clumps in the nuclear region of NGC 5135, $T_{d}$ could be warmer. In this
case, the calculated $M_{\rm mol}$ would be overestimated by roughly a factor
of $T_{d}$/(25 K). The resulting $M_{\rm mol}$, given in Column (11) of Table
2, ranges from $\sim$1 to 10 $\times 10^{8}\,M_{\odot}$.

To check the molecular gas mass, we derived the molecular gas mass from the CO
flux. We chose the conversion factor $\alpha_{\rm CO}=0.8$ M⊙(K km
s${}^{-1}\,$pc2)-1 (Downes & Solomon, 1998). The
${\mathrm{CO}}\,(6\text{--}5)$ /${\mathrm{CO}}\,(1\text{--}0)$ is about 2
(this will be explained in Sec. 4.3.1). The molecular gas mass derived this
way is consistent with the one derived from the dust continuum as shown in
Column (12) of Table 2.

Table 2 shows that all clumps have $M_{\rm mol}>M_{\rm vir}$ except for the
clumps b and d. The clumps in the former category (hereafter referred to as
Category (i)) are likely to be self-gravitationally bound or even undergoing
initial collapse. On the other hand, the clumps b and d in the other category
(hereafter Category (ii)) would require external pressure to remain bound.
Interestingly, the Category (ii) clumps are far away from on-going SF activity
and also show significantly higher CO (6-5)-to-dust continuum flux ratios (cf.
Table 2, Column (13)) than the clumps in Category (i).

We can also compare the clumps in NGC 5135 with the molecular gas clouds in
other galaxies. Fig. 8 is a plot of the cloud (FWHM) velocity dispersion as a
function of the cloud radius for the clumps in NGC 5135 as well as discrete
clouds in the center of the Milky Way, nearby spiral galaxies, and two
starburst galaxies NGC 253 and IC 5179. In comparison to the molecular clouds
over the disks of nearby normal galaxies, the clouds in the Milky Way center
observed by Oka et al. (2001) have a larger line width, but a smaller size. In
contrast, the clouds in the starburst galaxy NGC 253 (Leroy et al., 2015) and
the LIRG IC 5179 (Zhao et al., 2017) show both larger sizes and broader line
widths than clouds in the Milky Way center, but generally following the lines
of equal gas surface density for the case of Virialized clouds. The clouds in
the nuclear region of NGC 5135 are characterized by still larger sizes and
line widths.

Fig. 9 is a plot of the parameter, $\Delta V_{\rm FWHM}^{2}/R$, as a function
of molecular gas mass surface density $\Sigma$ for the same data set as in
Fig. 8. The thick diagonal line shows the locus of Virialized clouds. For
bound clouds clearly lying above this line, the cloud velocity width is likely
a manifestation of some external pressure. The dashed curves in Fig. 9, taken
from Field et al. (2011), indicate the relationship between $\Delta V^{2}_{\rm
FWHM}/R$ and $\Sigma$ for a varying external pressure. The Category (i) gas
clumps in NGC 5135 lie around the line tracing the Virial equilibrium. In
contrast, the two Category (ii) clumps are clearly located above the Virial
equilibrium, and require external pressure of the order of 108 cm-3K in order
to remain bound.

###  4.3. CO (6-5) Emission

####  4.3.1 CO (6-5) Emission to Continuum Ratio

On galaxy scale, the ratio of the ${\mathrm{CO}}\,(6\text{--}5)$ line
luminosity, $L_{\rm CO(6-5)}$, to $L_{\rm IR}$ varies only by up to 30% among
local LIRGs and shows little dependence on $L_{\rm IR}$ or the FIR color (Lu
et al., 2014, 2017). This strongly requires that the energy sources for both
the ${\mathrm{CO}}\,(6\text{--}5)$ and the dust emissions are ultimately tied
to the same SF process. This narrows down the candidate heating mechanisms for
the ${\mathrm{CO}}\,(6\text{--}5)$ emission to fewer choices, including photon
heating in the photon dominant regions (PDRs) around young massive stars and
SN-powered shock heating.

Our recent ALMA observation of nearby LIRGs show that, on scales of 100 pc or
less, local peaks of the ${\mathrm{CO}}\,(6\text{--}5)$ emission do not always
have corresponding peaks of the 435$\,\mu{\rm m}$ dust continuum emission.
Such examples include NGC 7130 (Zhao et al., 2016), IC 5179 (Zhao et al.,
2017) and the case of NGC 5135 shown here. Under the assumption of a constant
dust-to-gas abundance ratio, the spatial peaks of the two emissions should
follow each other if both the dust and ${\mathrm{CO}}\,(6\text{--}5)$
emissions are related to the same photon heating. This finding therefore
favors the SNe-powered shock gas heating scenario for the
${\mathrm{CO}}\,(6\text{--}5)$ emission. Here we investigate further this
subject in the case of NGC 5135.

As shown in Column (13) of Table 2, the ${\mathrm{CO}}\,(6\text{--}5)$ flux to
the 435$\,\mu{\rm m}$ continuum flux density ratio, $R_{\rm CO/cont}$, varies
among the CO (6-5) clumps. Further more, while the Category (i) clumps satisfy
$600\lesssim R_{\rm CO/cont}\lesssim 1800$ km s-1, the two Category (ii)
clumps have $R_{\rm CO/cont}>4,000$ km s-1. It is evident in Table 2, the
higher $R_{\rm CO/cont}$ values associated with the Category (ii) clumps are
mostly due to the unusually faint dust continuum emission at 435$\,\mu{\rm
m}$. One can express

$R_{\rm CO/cont}\propto(f_{\rm CO(6-5)}/f_{\rm CO(1-0)})\,(M_{\rm gas}/M_{\rm
dust})\,T^{-1}_{\rm d},$ (3)

where we have assumed the CO (1-0) flux, $f_{\rm CO(1-0)}$, scales with the
molecular gas mass $M_{\rm gas}$. This shows that a higher $R_{\rm CO/cont}$
can stem from either a hotter CO gas or/and a cooler dust temperature. In the
nuclear region of NGC 5135, the variation of $T_{d}$ is limited to, perhaps, a
factor of 3 (i.e., from 15 to 50 K) at most. The observed variation of $R_{\rm
CO/dust}$ is a factor of $\sim$5 in Table 2, mainly between the two clump
categories. A comparable variation is also seen in the case of IC 5179 (Zhao
et al. 2017). Therefore, it requires a modest variation of a factor of 2 or so
in $f_{\rm CO(6-5)}/f_{\rm CO(1-0)}$ in order to explain the observation. If
the ${\mathrm{CO}}\,(6\text{--}5)$ emission is associated with SNe-shock
heating, the ideal location for a higher $R_{\rm CO/cont}$ ratio is where
massive O star formation has ended while SNe activity is still strong, a
scenario we discussed in the case of NGC 7130 (Zhao et al. 2016). A necessary
condition for the validity of this scenario is that some dense gas can survive
the massive star formation, which might be possible in dense and clumpy ISM.

####  4.3.2 ${\mathrm{CO}}\,(6\text{--}5)$ Emission and Current SF

In Fig. 10, the black contours of the integrated
${\mathrm{CO}}\,(6\text{--}5)$ line emission are overlaid on a ground-based
8.7$\,\mu{\rm m}$ image (Díaz-Santos et al., 2008) on the left side, and on an
HST Pa-$\alpha$ image (Alonso-Herrero et al., 2006) on the right. In both
plots, we also show the low-resolution Chandra 0.4-8 keV broadband X-ray
emission (Levenson et al. 2004) in red contours. This X-ray emission is mostly
associated with a hot, ionized gas powered by SNRs (Levenson et al 2004;
Colina et al. 2012).

The 8.7$\,\mu{\rm m}$ image is dominated by the mid-IR emission bands from the
so-called Polycyclic Aromatic Hydrocarbons (PAHs). The Pa-$\alpha$ emission is
caused by the ionizing UV radiation from massive O stars and is regarded as a
reasonable tracer of the on-going SF activity (timescale $\sim$ 10 Myr). The
PAH emission traces SF of 10 times longer time scale ($\sim$100 Myr), as PAH
molecules are mostly heated by non-ionizing B stars (Díaz-Santos et al.,
2008). As shown in Fig. 10, the PAH and Pa-$\alpha$ emissions show similar
surface brightness distributions. In contrast, the overall spatial morphology
of the ${\mathrm{CO}}\,(6\text{--}5)$ emission appears to be different from
that of either the PAH or the Pa-$\alpha$ emission. However, the Category (i)
${\mathrm{CO}}\,(6\text{--}5)$ clumps are all relatively close to local
emission peaks of the PAH or Pa-$\alpha$ emission whereas the Category (ii)
${\mathrm{CO}}\,(6\text{--}5)$ clumps are significantly farther away from any
bright PAH or Pa-$\alpha$ peak. Therefore, it is reasonable to expect a much
weaker far-UV radiation intensity at the location of each Category (ii) clump.
This naturally explains why the dust emission is unusually faint at each of
the Category (ii) clumps.

###  4.4. Possible Heating Scenarios for ${\mathrm{CO}}\,(6\text{--}5)$
Emission

####  4.4.1 SN-powered Shock Heating Scenario

With an integral field spectrograph, Colina et al. (2012) measured the
intensity and velocity fields of both the Br-$\gamma$ and the [Fe ii]
1.64$\,\mu{\rm m}$ emission lines in the nuclear region of NGC 5135. While the
Br-$\gamma$ traces the current star formation, the [Fe ii] line emission is
regarded as a particularly good tracer of SNRs (Greenhouse et al., 1991). In
Fig. 11, we show a plot of the [Fe ii]-to-Br-$\gamma$ line ratio versus the CO
(6-5)-to-continuum flux density ratio for all the clumps listed in Table 2,
except for clump d, which is located outside the field of view of the [Fe ii]
observation. We also indicated the typical [Fe ii]-to-Br-$\gamma$ line ratios
for different astrophysical objects, taken from Falcón-Barroso et al. (2014).
Note that the line ratio range shown for Seferts is largely irrelevant here as
our molecular clouds are all located far away from the AGN.

A number of studies have attempted to identify the physical causes behind the
observed variations of the [Fe ii]-to-Br-$\gamma$ line ratio (e.g., Alonso-
Herrero et al. 1997; Moorwood & Oliva 1988; Mouri et al. 1990, 1993;
Greenhouse et al. 1991; Rodríguez-Ardila et al. 2004, 2005; Ramos Almeida et
al. 2006, 2009; Riffel et al. 2013; Falcón-Barroso et al. 2014). From these
studies, two main conditions for an enhanced [Fe ii] emission relative to a
Hydrogen recommbination line emission emerge: (a) Presence of shocked gas and
(b) favorable environment for the gas-phase Fe to be abudant in the form of
Fe+. Iron is normally depleted onto grains in the interstellar gas phase, so
fast shocks, such as those associated with SNe, can cause grain destructions
and therefore enrich gas-phase Fe abundance. Another important prerequisite
for a strong [Fe ii] line is an ionization field in favor of Fe+. Given the
low ionization potential of Fe+ (16.2 eV), most of Fe is in higher ionization
states in HII regions. In comparison, partially ionized gas in SNRs and Fe+ is
believed to be abundant (Moorwood & Oliva, 1988). The collisional exitation
with electrons could therefore make the [Fe ii] line much brighter in SNRs.

The [Fe ii]-to-Br-$\gamma$ line ratios for the Category (i) clumps in NGC 5135
are around 3, which is just outside the upper tip of the range for Galactic
HII regions. These ratios are slightly higher than those seen in the nuclear
star-forming regions in the Seyfert galaxy NGC 613 (Falcón-Barroso et al.,
2014), suggesting some mild enhancement in the [Fe ii] line emission for the
gas clouds in NGC 5135. In comparison, this line ratio for the Category (ii)
clump b equals 6.5, implying a factor of 2 further enhancement in the relative
[Fe ii] emission from the Category (i) clouds. It is not surprising for the
observed line ratio of the Category (ii) cloud to be smaller than the typical
values seen in Galactic SNRs because we are averaging over a much larger area
than the typical size of Galactic SNRs and also because there is still low-
surface brightness star forming activity near the cloud (see Fig. 10). The
observed trend in Fig. 11 indicates that the same SNe shocks are likely to
play a positive role in the observed variations in both [Fe ii]/Br-$\gamma$
and CO (6-5)/dust ratios.

Additional evidence in favor of the SNe-shock heating scenario for the
${\mathrm{CO}}\,(6\text{--}5)$ emission in NGC 5135 includes (a) the
prevailing X-ray emission from the hot, ionized gas excited via SN shocks
(Colina et al. 2012); and (b) that there is a very good velocity field
correspondence between the ${\mathrm{CO}}\,(6\text{--}5)$ clumps and that of
the underlying [Fe ii] emission: the nominal ${\mathrm{CO}}\,(6\text{--}5)$
and [Fe ii] line velocity offset varies around the mean of -142 km s-1 by only
a few  km s-1 among the clumps. (Note that the mean velocity field difference
is likely a result of the different velocity reference frames adopted.) This
velocity correspondence suggests that the warm CO gas and the shocked/ionized
gas are reasonably well mixed with each other in space and velocity field.
Another independent evidence in favor of the SNe-shock heating scenario is the
global tight correlation between the IR dust emission and the mid-$J$ CO line
emission shown by Lu et al. (2017), which requires that the gas heating
ultimately derives the energy from the same SF process. The SNe heating
scenario would naturally fit this requirement.

####  4.4.2 Bar-Induced Shock Heating Scenario

Fig. 12 displays the integrated ${\mathrm{CO}}\,(6\text{--}5)$ line emission
contours overlaid on the HST F606W (0.606$\,\mu{\rm m}$) image (Malkan et al.,
1998) and HST F160W (1.60$\,\mu{\rm m}$) image (Alonso-Herrero et al., 2006),
respectively. The HST images aligned with our CO (6-5) data by matching our
adopted AGN position with the brightest point in each optical image. As
already mentioned before, the ${\mathrm{CO}}\,(6\text{--}5)$ emission has a
good spatial correspondence with the dust lanes that can be seen in the
optical and near-IR continuum images here. These roughly symmetrical dust
lanes are induced by the inner stellar bar, both of which are more visible in
a larger UV/optical image such as the one shown by Mulchaey & Regan (1997).
This correspondence between the ${\mathrm{CO}}\,(6\text{--}5)$ emission and
the bar-induced dust lanes in NGC 5135 is similar to that observed in NGC
7130, another LIRG with a strong stellar bar (Zhao et al., 2016). Indeed NGC
5135 and NGC 7130 have many similarities: both are LIRGs with a strong
circumnuclear star formation and a Seyfert 2 nucleus. Circumnuclear dust lanes
have been found in many spiral galaxies, though strong 2-arm dust lanes are
found only in barred galaxies such as NGC 5135 and NGC 7130 (Martini et al.,
2003).

According to the model of Athanassoula (1992), the 2-arm dust lanes are
associated with the shock fronts triggered by the presence of a bar in a
rotating gas disk. Thus, bar-induced shocks could be possible in the nuclear
region of NGC 5135. Inside the dust lanes, the gas (and dust) density is
significantly enhanced, but SF is suppressed by strong shears (Athanassoula,
1992). This seems to imply that the warm dense gas traced by
${\mathrm{CO}}\,(6\text{--}5)$ and the SF regions traced by Pa-$\alpha$ are
not related to each other.

In order to account for the tight correlation between the
${\mathrm{CO}}\,(6\text{--}5)$ emission and the total dust emission on galaxy
scale for LIRGs, one has to relate the ${\mathrm{CO}}\,(6\text{--}5)$ emitting
gas to the SF activity in this scenario. It is still possible that the warm
dense gas and the SF regions are related to each other, albeit their positions
are slightly offset. It is known that some galaxies with weak stellar bars
(therefore weaker shears in dust lanes) have SF in their dust lanes (Comte &
Duquennoy, 1982; Martini et al., 2003). Hypothetically, one can envisage the
following scenario: First, SF does occur in clouds of dense gas formed in the
post-shock gas down stream from the bar-induced shock front (the dust lane).
Then these dense gas clouds will be rapidly consumed/destroyed by the SF and
the associated feedback. In this scenario, under the assumption that the
destruction time scale of the dense clouds is much shorter than the SF time
scale associated with the Pa-$\alpha$ emission (a few Myrs), the spatial
offset is the product of the SF time scale times the downstream velocity of
the post-shock gas, which is a few 10s  km s-1(Athanassoula, 1992). This
indeed results in an estimate for the offset of $\lesssim$100 pc. It is worth
noting that similar offset between HII regions and dust lanes associated with
spiral arms in grand-design galaxies such as M 51 have been found in the
literature, and Scoville et al. (2001) argued that it implies that the HII
regions develop subsequent to the time of maximum concentration of the dust
and molecular clouds.

However, it is unclear how this bar-induced shock heating scenario for the CO
(6-5) emission can be made to explain the similar variation in the CO
(6-5)-to-continuum flux ratio seen in IC 5179 (Zhao et al. 2017), which does
not have a strong stellar bar.

##  5\. SUMMARY

In this paper we present the results from our ALMA observations of the
${\mathrm{CO}}\,(6\text{--}5)$ line and its underlying dust continuum at
435$\,\mu{\rm m}$ in the nuclear region of the nearby LIRG, Seyfert 2 galaxy
NGC 5135, at a physical resolution of 48 $\times$ 40 pc. Our main findings
are:

(1) The central AGN is undetected in either the 435 $\,\mu{\rm m}$ dust
continuum or ${\mathrm{CO}}\,(6\text{--}5)$ line emission if its line velocity
width is no less than $\sim$40  km s-1, resulting in an AGN contributes at
most 1% of the integrated circum-nuclear ${\mathrm{CO}}\,(6\text{--}5)$ flux
seen in our ALMA observation. On the other hand, the non-detection in
continuum emission may simply reflect the lack of sensitivity in our
observation.

(2) The circum-nuclear ${\mathrm{CO}}\,(6\text{--}5)$ emission is resolved
into gas clumps of radii of 45-180 pc and line velocity widths of 60-88  km
s-1. While the clump sizes are only slightly larger than typical giant
molecular clouds in nearby spiral galaxies, their velocity widths are
significantly higher. They fall into two categories: (i) The 5 clumps that are
near some current star formation activity are likely to be in Virial
equilibrium, and (ii) the other 2 clumps without clear current star formation
activity nearby seem to be unbound unless there is significant external
pressure.

(3) The clumps in Category (ii) have much higher
${\mathrm{CO}}\,(6\text{--}5)$ to dust continuum ratios than those in Category
(i). Furthermore, the CO (6-5)-to-continuum ratios show an increasing trend
with the [Fe ii]-to-Br-$\gamma$ ratios, which we interpret as evidence for
supernova-driven shocked gas providing a significant contribution to the
${\mathrm{CO}}\,(6\text{--}5)$ emission.

(4) The clumps are distributed along the symmetric optical dust lanes
associated with the stellar bar at the center of the galaxy. Like NGC 7130,
another barred Seyfert galaxy, the gas concentrations could be a result of the
bar-induced instability and is subject to bar-induced shock heating.

## Acknowledgments

We thank an anonymous referee for a number of very constructive comments. We
thank Drs. Cheng Cheng, Luis Colina, Adam Leroy, Claudio Ricci, and Chentao
Yang for their insightful comments and/or useful communications during the
preparation of the manuscript. This paper makes use of the following ALMA
data: ADS/JAO.ALMA#2013.1.00524.S. ALMA is a partnership of ESO (representing
its member states), NSF (USA), and NINS (Japan), together with NRC (Canada)
and NSC and ASIAA (Taiwan), in cooperation with the Republic of Chile. The
Joint ALMA Observatory is operated by ESO, AUI/NRAO, and NAOJ. This work is
supported in part by the National Key R&D Program of China grant
#2017YFA0402704, the NSFC grant #11673028 and #11673057, and by the Chinese
Academy of Sciences (CAS), through a grant to the CAS South America Center for
Astronomy (CASSACA) in Santiago, Chile. C.C. acknowledges support by NSFC
grant No. 11503013. Y.G. acknowledges support by NSFC grants No. 11173059,
11390373, and 11420101002. H.W. acknowledges support by NSFC grant No.
11733006. V.K. acknowledges support from the FONDECYT grant No. 3160117.
T.D.-S. acknowledges support from ALMA-CONICYT project 31130005 and FONDECYT
regular project 1151239.

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Figure 1.—  Panels (a) and (b): Contours of the frequency-integrated CO (6-5)
intensity overlaid on the images of the same integrated CO (6-5) intensity in
(a) and the 435$\,\mu{\rm m}$ dust continuum emission in (b). Panels (c) and
(d) are respectively the line velocity field (moment 1) and the velocity
dispersion (moment 2) maps of the CO (6-5) line emission obtained from the uv-
taper image. The images in (c) and (d) are generated by using only those
spaxels above 4-$\sigma_{ch}$, where $\sigma_{ch}$ is the r.m.s noise per
frequency channel ($\sigma_{ch}$ = 40 mJy beam-1 for uv-taper image). The
contours in panels (a) and (b) are shown at [3,5]$\times\sigma$ (where the
noise $\sigma$ = 1.2 Jy beam-1 kms-1). The unit of the color bar in each panel
is given near the upper-right color. The filed ellipse in white near the
lower-left corner in (a) or (b) is the ALMA beam. The large ellipses in red in
(a) mark the regions for spectrum and flux extractions given in Table 2. The
red plus sign marks the AGN position adopted. The figures are before the
primary beam correction integrated over the barycentric velocities from 3,971
to 4,157 km s-1.  Figure 2.—  Channel maps of the CO (6-5) line emission (in
contours), each overlaid on the image of the total, frequency-integrated CO
(6-5) emission (e.g., from Fig. 1a). The channel interval is 13.5 km s-1, with
the channel central (barycentric) velocity shown in each channel map.  Figure
3.— Same as the panels (a) and (b) in Fig. 1, respectively, but using the CO
(6-5) and continuum data at a larger effective beam. The contour levels are
[3, 4, 5, 6]$\times\sigma$ ($\sigma=8$ Jy beam-1 km s-1, with the beam size of
0.4″ $\times$ 0.4″ here as shown by the filled ellipse in white in each panel.
Figure 4.— Black contours of the integrated CO (6-5) line emission overlaid on
(a) an VLA radio cyan contours at 6cm (the cyan contour levels are [3, 5, 7,
10, 30, 50]$\times\sigma$ ($\sigma$ =1.2e-04 Jybeam-1) and (b) a Chandra 4-8
kev X-ray image. The black contour level are [3, 5, 6]$\times\sigma$
($\sigma$=1.2 Jybeam-1kms-1). The white plus sign in each panel presents the
adopted AGN location.  Figure 5.—  (a) Contours of the integrated CO (6-5)
from 4019 to 4032 km/s on the image of the same integrated CO (6-5) intensity.
The cyan arrow points to the central AGN. The contour levels are [3, 5,
8]$\times\sigma$ ($\sigma$=0.5 Jybeam-1kms-1). And (b) the CO (6-5) spectrum
at the central AGN. The central velocity ($V_{0}$) and FWHM of a Gaussian fit
are given in the plot.  Figure 6.—  Plot of the empirical infrared spectrum
(the green curve) of the AGN in NGC 5135, which is based the observed X-ray
luminosity and anchored at the 12$\,\mu{\rm m}$ flux density measurement (the
filled circle in blue). Also shown are two ALMA flux upper limits at
435$\,\mu{\rm m}$ and 1.3 mm, respectively (see the text).  Figure 7.—
Spatially integrated CO (6-5) line profile of various clumps (a1, a2, a3, a4,
a5, b, c, d). The central velocity ($V_{0}$) and FWHM of a Gaussian fit are
given in each plot.  Figure 8.— Plot of the CO (6-5) line width, $\Delta
V_{\rm FWHM}$, as a function of the cloud radius $R$ for a sample of giant
molecular clouds and gas complexes in the Milky Way and various local
galaxies, adopted from Leroy et al. (2015). As a comparison, our NGC 5135
clumps are added (i.e., large filled circles labeled by the clump number). The
light dashed lines follow $\Delta V_{\rm FWHM}\propto$ R0.5, the relation
expected for Virialized clouds with a fixed surface density $\Sigma$, spaced
by a factor of 2 vertically. The thick dashed line is for $\Sigma$ $\approx$
285 M$\odot$ pc-2.  Figure 9.—  Plot of the cloud radius-line width
coefficient, $\Delta V^{2}_{\rm FWHM}/R$, as a function of the gas surface
density $\Sigma$, also adopted from Leroy et al. (2015), for the same data set
as in Fig. 7. The gas clouds that are in Virial equilibrium (i.e., $M_{\rm
mol}=M_{\rm vir}$) follow the thick line. Bound clouds with $M_{\rm
mol}<M_{\rm vir}$ follow one of the thin curves representing various external
pressures as labelled in terms of P/kB (in units of cm-3 K).  Figure 10.—
Integrated CO (6-5) line emission contours, at [3, 5, 6]$\times\sigma$ (where
$\sigma$=1.2 Jybeam-1kms-1), overlaid on (a) a 8.7$\,\mu{\rm m}$ image
dominated the PAH emission (in log scale) and (b) an image of the Pa-$\alpha$
line emission (in log scale). The red contours in each panel stands for the
X-ray intensity of NGC 5135 (with the contours at 5, 20, 30, 40, 70, 150, 260
counts), obtained in the Chandra 0.4-8 keV band by Levenson et al. (2004). The
white plus sign in each panel marks the adopted AGN location.  Figure 11.—
Plot of the [Fe ii] 1.46$\,\mu{\rm m}$-to-Br-$\gamma$ line ratio as a function
of the CO (6-5)-to-continuum flux ratio for the NGC 5135 clumps as labelled.
The extinction corrected [Fe ii] and Br-$\gamma$ line surface brightnesses at
the local of a clump are estimated from Bedregal et al. (2009) and the typical
error is 0.15 dex. The typical ratios for different types of astrophysical
objects are noted in the plot (see the text). The clump d is not plotted here
as it does not have the corresponding [Fe ii] or Br-$\gamma$ data.  Figure
12.—  Integrated CO (6-5) line emission contours, at [3, 5, 6]$\times\sigma$
(where $\sigma$=1.2 Jybeam-1kms-1), overlaid on (a) an HST F606W
(0.606$\,\mu{\rm m}$, in log scale) image and (b) an HST F160W (1.60$\,\mu{\rm
m}$, in log scale) image. The white plus sign in each panel marks the AGN
position.

Table 1Basic properties of NGC 5135 and ALMA Observation Log |  |  | Basic Properties |  |  |  |
---|---|---|---|---|---|---|---
Name | R.A. (J2000) | Dec. (J2000) | Dist. | cz | Morph | Spectral Type | $\log\,L_{\rm IR}$
| (hh:mm:ss) | (dd:mm:ss) | Mpc | kms-1 |  |  | $L_{\odot}$
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8)
NGC 5135 | 13:25:43.99 | -29:50:01.06 | 59 | 4105 | SB (s)ab | Sy 2 | 11.33
|  |  | ALMA observation log |  |  |  |
SB | Date | Time (UTC) | Configuration | Nant | lmax | tint | Tsys
| (yyyy/mm/dd) |  |  |  | (m) | (seconds) | (K)
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8)
Xa216e2$\\_$Xcb0 | 2015/06/02-2015/06/03 | 23:58:10-00:51:44 | C34-5 | 39 | 885.6 | 21.03 | 935 - 839

Table 2 Physical properties of the individual clumps in our CO (6-5) image

No. | Size | PA | Radius | $V_{0}$ | $\Delta V_{\rm FWHM}$ | fpeak | fCO(6-5) | Icont | Mvir | Mmol | M${}^{*}_{\rm mol}$ | RCO/cont
---|---|---|---|---|---|---|---|---|---|---|---|---
|  ′′$\times$′′ | Deg | pc | kms-1 | kms-1 | Jy beam-1 | Jy kms-1 | mJy |  $\times$107 M$\odot$ |  $\times$10${}^{8}\,$M$\odot$ |  $\times$10${}^{8}\,$M$\odot$ | kms-1
|  |  |  |  |  |  km s-1 |  |  |  |  |  |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10) | (11) | (12) | (13)
a1 | 0.65$\times$0.34 | 16.50 | 99.2 | 4062.9 | 88.0 | 13.3 | 91.1 | 81.8 | 37.0 | 10.5 | 12.6 | 1113.0
|  $\pm$0.019$\times$0.010 |  |  $\pm$2.2 |  $\pm$3.0 |  $\pm$7.6 |  |  $\pm$3.0 |  $\pm$8.4 |  $\pm$7.4 |  $\pm$2.9 |  $\pm$0.4 |  $\pm$119.0
a2 | 0.36$\times$0.10 | -0.45 | 44.1 | 4084.4 | 63.6 | 9.2 | 15.5 | 11.4 | 7.6 | 1.4 | 2.1 | 1362.0
|  $\pm$0.025$\times$0.012 |  |  $\pm$3.0 |  $\pm$3.7 |  $\pm$9.2 |  |  $\pm$3.1 |  $\pm$2.7 |  $\pm$2.7 |  $\pm$0.5 |  $\pm$0.4 |  $\pm$419.0
a3 | 0.40$\times$0.25 | 0.74 | 62.6 | 4053.7 | 59.1 | 10.6 | 29.6 | 28.7 | 8.9 | 3.7 | 4.1 | 1032.0
|  $\pm$0.023$\times$0.014 |  |  $\pm$2.9 |  $\pm$4.0 |  $\pm$6.0 |  |  $\pm$2.1 |  $\pm$6.2 |  $\pm$2.3 |  $\pm$1.2 |  $\pm$ 0.3 |  $\pm$236.0
a4 | 0.64$\times$0.19 | 0.13 | 54.1 | 4047.6 | 62.6 | 9.6 | 28.5 | 21.6 | 9.4 | 2.6 | 3.9 | 1321.0
|  $\pm$0.054$\times$0.011 |  |  $\pm$4.0 |  $\pm$2.4 |  $\pm$7.2 |  |  $\pm$1.8 |  $\pm$2.4 |  $\pm$2.7 |  $\pm$0.7 |  $\pm$0.2 |  $\pm$171.0
a5 | 0.72$\times$0.44 | 3.38 | 123.9 | 4027.6 | 61.7 | 9.6 | 55.0 | 39.7 | 19.7 | 5.1 | 7.6 | 1850.0
|  $\pm$0.023$\times$0.022 |  |  $\pm$4.3 |  $\pm$1.2 |  $\pm$5.6 |  |  $\pm$1.5 |  $\pm$6.3 |  $\pm$4.5 |  $\pm$1.5 |  $\pm$0.2 |  $\pm$221.0
b | 0.52$\times$0.38 | 5.09 | 95.7 | 4089.6 | 76.7 | 8.9 | 55.3 | 13.5 | 25.9 | 1.7 | 7.6 | 4084.0
|  $\pm$0.028$\times$0.026 |  |  $\pm$4.4 |  $\pm$1.2 |  $\pm$4.4 |  |  $\pm$2.0 |  $\pm$1.9 |  $\pm$3.6 |  $\pm$0.5 |  $\pm$0.3 |  $\pm$598.0
c | 0.72$\times$0.21 | -32.46 | 76.1 | 4027.8 | 69.1 | 9.4 | 41.8 | 48.3 | 16.0 | 6.2 | 5.7 | 866.0
|  $\pm$0.025$\times$0.007 |  |  $\pm$2.1 |  $\pm$1.9 |  $\pm$5.0 |  |  $\pm$2.8 |  $\pm$5.9 |  $\pm$2.8 |  $\pm$1.7 |  $\pm$0.4 |  $\pm$120.0
daafootnotemark:  | 0.94$\times$0.84 | 0.00 | 182.0 | 4067.3 | 87.3 | 30.2 | 69.4 | 13.0 | 66.6 | 1.7 | 9.6 | 5320.0
|  $\pm$0.016$\times$0.014 |  |  $\pm$2.4 |  $\pm$1.7 |  $\pm$4.8 |  |  $\pm$5.5 |  $\pm$3.8 |  $\pm$8.4 |  $\pm$0.7 |  $\pm$0.8 |  $\pm$1603.0

11footnotetext: Using the image in Fig. 3d, corresponding to a larger beam of
0.4″ $\times 0.4$″.