# Diverse Molecular Structures Across The Whole Star-Forming Disk of M83:
High fidelity Imaging at 40pc Resolution

Jin Koda  Department of Physics and Astronomy, Stony Brook University, Stony
Brook, NY 11794-3800, USA  Akihiko Hirota  NAOJ Chile, National Astronomical
Observatory of Japan, Los Abedules 3085 Office 701, Vitacura, Santiago
763-0414, Chile  Joint ALMA Observatory, Alonso de Córdova 3107, Vitacura,
Santiago 763-0355, Chile  Fumi Egusa  Institute of Astronomy, Graduate School
of Science, The University of Tokyo, 2-21-1 Osawa, Mitaka, Tokyo 181-0015,
Japan  Kazushi Sakamoto  Academia Sinica, Institute of Astronomy and
Astrophysics, Taipei 10617, Taiwan  Tsuyoshi Sawada  NAOJ Chile, National
Astronomical Observatory of Japan, Los Abedules 3085 Office 701, Vitacura,
Santiago 763-0414, Chile  Joint ALMA Observatory, Alonso de Córdova 3107,
Vitacura, Santiago 763-0355, Chile  Mark Heyer  Department of Astronomy,
University of Massachusetts Amherst, 710 North Pleasant Street, Amherst, MA
01003, USA  Junichi Baba  National Astronomical Observatory of Japan, 2-21-1
Osawa, Mitaka, Tokyo 181-8588, Japan.  Department of Astronomy, School of
Science, Graduate University for Advanced Studies (SOKENDAI), 2-21-1 Osawa,
Mitaka, Tokyo, 181-1855 Japan  Samuel Boissier  Aix Marseille Univ., CNRS,
CNES, Laboratoire d’Astrophysique de Marseille, Marseille, France  Daniela
Calzetti  Department of Astronomy, University of Massachusetts Amherst, 710
North Pleasant Street, Amherst, MA 01003, USA  Jennifer Donovan Meyer
National Radio Astronomy Observatory, 520 Edgemont Rd, Charlottesville, VA
22903, USA  Bruce G. Elmegreen  IBM Research Division, T. J. Watson Research
Center, 1101 Kitchawan Road, Yorktown Heights, NY 10598, USA  Armando Gil de
Paz  Departamento de Física de la Tierra y Astrofísica, Facultad de CC.
Físicas, Universidad Complutense de Madrid, 28040, Madrid, Spain  Instituto de
Física de Partículas y del Cosmos IPARCOS, Facultad de CC. Físicas,
Universidad Complutense de Madrid, 28040 Madrid, Spain  Nanase Harada
National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo
181-8588, Japan.  Department of Astronomy, School of Science, Graduate
University for Advanced Studies (SOKENDAI), 2-21-1 Osawa, Mitaka, Tokyo,
181-1855 Japan  Luis C. Ho  Kavli Institute for Astronomy and Astrophysics,
Peking University, Beijing 100871, China  Department of Astronomy, School of
Physics, Peking University, Beijing 100871, China  Masato I.N. Kobayashi
National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo
181-8588, Japan.  Nario Kuno  Division of Physics, Faculty of Pure and Applied
Sciences, University of Tsukuba, 1-1-1 Tennodai, Tsukuba, Ibaraki 305-8577,
Japan  Tomonaga Center for the History of the Universe (TCHoU), University of
Tsukuba, 1-1-1 Tennodai, Tsukuba, Ibaraki 305-8577, Japan  Amanda M Lee
Department of Physics and Astronomy, Stony Brook University, Stony Brook, NY
11794-3800, USA  Department of Astronomy, University of Massachusetts Amherst,
710 North Pleasant Street, Amherst, MA 01003, USA  Barry F. Madore  The
Observatories, Carnegie Institution for Science, 813 Santa Barbara Street,
Pasadena CA 91101 USA  Fumiya Maeda  Institute of Astronomy, Graduate School
of Science, The University of Tokyo, 2-21-1 Osawa, Mitaka, Tokyo 181-0015,
Japan  Sergio Martín  European Southern Observatory, Alonso de Córdova, 3107,
Vitacura, Santiago 763-0355, Chile  Joint ALMA Observatory, Alonso de Córdova
3107, Vitacura, Santiago 763-0355, Chile  Kazuyuki Muraoka  Department of
Physics, Graduate School of Science, Osaka Metropolitan University, 1-1
Gakuen-cho, Naka-ku, Sakai, Osaka 599-8531, Japan  Kouichiro Nakanishi
National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan
Department of Astronomy, School of Science, Graduate University for Advanced
Studies (SOKENDAI), 2-21-1 Osawa, Mitaka, Tokyo, 181-1855 Japan  Sachiko
Onodera  School of Science and Engineering, Meisei University, Hino, Tokyo
191-8506, Japan  Jorge L. Pineda  Jet Propulsion Laboratory, California
Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109-8099, USA
Nick Scoville  Division of Physics and Astronomy, California Institute of
Technology, Pasadena, CA 91125, USA  Yoshimasa Watanabe  Materials Science and
Engineering, College of Engineering, Shibaura Institute of Technology, 3-7-5
Toyosu, Koto-ku, Tokyo 135-8548, Japan

###### Abstract

We present ALMA imaging of molecular gas across the full star-forming disk of
the barred spiral galaxy M83 in CO($J$=1-0). We jointly deconvolve the data
from ALMA’s 12m, 7m, and Total Power arrays using the MIRIAD package. The data
have a mass sensitivity and resolution of $10^{4}~{}{\rm M_{\odot}}$
($3\sigma$) and 40 pc – sufficient to detect and resolve a typical molecular
cloud in the Milky Way with a mass and diameter of $4\times 10^{5}~{}{\rm
M_{\odot}}$ and $40\,{\rm pc}$, respectively. The full disk coverage shows
that the characteristics of molecular gas change radially from the center to
outer disk, with the locally measured brightness temperature, velocity
dispersion, and integrated intensity (surface density) decreasing outward. The
molecular gas distribution shows coherent large-scale structures in the inner
part, including the central concentration, offset ridges along the bar, and
prominent molecular spiral arms. However, while the arms are still present in
the outer disk, they appear less spatially coherent, and even flocculent.
Massive filamentary gas concentrations are abundant even in the interarm
regions. Building up these structures in the interarm regions would require a
very long time ($\gtrsim 100$ Myr). Instead, they must have formed within
stellar spiral arms and been released into the interarm regions. For such
structures to survive through the dynamical processes, the lifetimes of these
structures and their constituent molecules and molecular clouds must be long
($\gtrsim 100\,{\rm Myr}$). These interarm structures host little or no star
formation traced by H$\alpha$. The new map also shows extended CO emission,
which likely represents an ensemble of unresolved molecular clouds.

galaxies: individual (M83, NGC5236) — galaxies: ISM — galaxies: spiral

††facilities: ALMA††software: CASA, MIRIAD, DS9, BBarolo

##  1 Introduction

Molecular gas is a key link between star formation and galaxy evolution across
cosmic time (Tacconi et al., 2020). In particular, molecular clouds
($10^{4-7}~{}{\rm M_{\odot}}$) host virtually all star formation in the local
Universe, and their formation and evolution are the first crucial steps
leading to star formation. Galactic dynamics around stellar bars and spiral
arms stir the gas and stimulate the formation and evolution of molecular
clouds. Subsequent star formation and feedback into the gas are two of the
most important factors in galaxy growth and quenching, respectively. These
processes leave imprints on the distribution and physical conditions of
molecular gas and clouds over a galactic disk. To gain a more complete
understanding of these processes, it is essential to map the full population
of molecular gas and clouds over a galaxy.

Until recently, molecular gas studies over large spiral galaxies have been
limited by sensitivity and resolution. The previous CO(1-0) studies of M51 had
a cloud detection limit of $\gtrsim 10^{5}~{}{\rm M_{\odot}}$, and left
significant CO emission undetected or unresolved below this limit (Koda et
al., 2009, 2011; Schinnerer et al., 2013; Pety et al., 2013). The molecular
clouds with mass $\gtrsim 10^{5}~{}{\rm M_{\odot}}$ in the Milky Way (MW)
contain only $\sim$half of the total molecular mass (Scoville & Sanders, 1987;
Heyer & Dame, 2015), raising a natural question as to how the other hidden
half of molecular gas, presumably within smaller clouds with $\lesssim
10^{5}~{}{\rm M_{\odot}}$, contributes to gas evolution and star formation. A
cloud-scale resolution of $40\,{\rm pc}$, the typical diameter of Galactic
molecular clouds (Scoville & Sanders, 1987), is necessary for mapping the full
cloud population and resolving cloud properties (Donovan Meyer et al., 2012;
Schinnerer et al., 2013; Donovan Meyer et al., 2013; Colombo et al., 2014;
Rosolowsky et al., 2021).

Even with ALMA, deep observations covering large galactic disks are expensive
in time. The recent ALMA large survey of nearby galaxies (PHANGS; Leroy et
al., 2021a) set the target sensitivity to $\sim 10^{5}~{}{\rm M_{\odot}}$ and
the resolution to $\sim 100\,{\rm pc}$, and observed only the inner parts of
galactic disks. In addition, many large surveys with ALMA rely on the excited
transition of CO(2-1), instead of CO(1-0), for a higher observational
efficiency (by a factor of $\sim 8$). This compromise is inevitable for large
surveys, but could introduce a bias toward warmer and/or denser gas. Indeed,
environmental variations of the CO 2-1/1-0 line ratio have been detected in
galaxies (Sakamoto et al., 1997; Koda et al., 2012, 2020; Yajima et al., 2021;
den Brok et al., 2021). We therefore need CO(1-0) observations to take a full
census of molecular gas in galactic disks.

Currently, the only extragalactic studies of lower-mass clouds ($\sim
10^{4}~{}{\rm M_{\odot}}$) throughout a disk are limited to the closest dwarfs
and dwarf-like spirals (LMC, SMC, M33, and NGC300; e.g., Mizuno et al., 2001;
Fukui et al., 2008; Rosolowsky et al., 2003; Tosaki et al., 2011; Gratier et
al., 2012; Faesi et al., 2014). However, the gas evolution in these galaxies
appears substantially different from that in large spiral galaxies, such as
the MW (Koda et al., 2016). Much deeper observations of a substantial spiral
galaxy, even a single one, provide an unbiased picture of molecular gas and
cloud evolution over a large galactic disk.

We present deep mosaic observations of the nearby barred spiral galaxy M83
(NGC 5236) in the fundamental CO(1-0) transition with ALMA, which provides a
mass sensitivity and spatial resolution of $10^{4}~{}{\rm M_{\odot}}$
($\gtrsim 3\sigma$) and $40~{}\,{\rm pc}$ respectively. We use the ALMA 12m,
7m, and Total Power (TP) arrays, convert the TP image into visibilities and
jointly-deconvolve them for high quality images (Koda et al., 2019). This
paper presents the first results on the distribution of molecular gas and
clouds and their relation to star formation in M83. Subsequent papers will
present a more in-depth analysis on, e.g., cloud properties and physical
conditions as well as global and local environments of star formation.

This CO(1-0) study is complementary to large ALMA surveys of galaxies,
including PHANGS (Sun et al., 2018; Schinnerer et al., 2019; Leroy et al.,
2021a), ALCHEMI (Martín et al., 2021; Harada et al., 2019, 2022), VERTICO
(Brown et al., 2021; Zabel et al., 2022), and ALMA JERRY (Jachym in private
communication), as well as many other studies of individual galaxies.

###  1.1 The Barred Spiral Galaxy M83

M83 is one of the closest barred spiral galaxies at the distance of
$d=4.5\,{\rm Mpc}$ (Thim et al., 2003). It closely resembles our own MW
(Churchwell et al., 2009), and its nearly face-on geometry, with an
inclination angle of $i=26\deg$ (this study) shows the full structure of the
disk without the distance ambiguity that MW observations encounter. Table 1
summarizes the parameters of M83. The galaxy’s total stellar mass of $M_{\rm
star}=2.5\times 10^{10}~{}{\rm M_{\odot}}$ and exponential disk scale length
of $h=1.7\,{\rm kpc}$ (Barnes et al., 2014), assuming a distance of 4.5 Mpc,
are slightly smaller than those of the MW ($4.6\times 10^{10}~{}{\rm
M_{\odot}}$ and $2.2\,{\rm kpc}$; Bovy & Rix, 2013). The molecular gas mass of
$M_{\rm H_{2}}=2.6\times 10^{9}~{}{\rm M_{\odot}}$ (this study) and star
formation rate of ${\rm SFR}=5.2\,~{}{\rm M_{\odot}}/\,{\rm yr}$ (Jarrett et
al., 2019) are slightly higher than those of the MW ($1.0\times 10^{9}~{}{\rm
M_{\odot}}$ and $1.9\,~{}{\rm M_{\odot}}/\,{\rm yr}$; Heyer & Dame, 2015;
Chomiuk & Povich, 2011). Therefore, M83 is only slightly smaller and slightly
more active in star formation than the MW. The metallicity of M83 is
12+log(O/H)=$8.78\pm 0.07$ in stars and $8.73\pm 0.27$ in gas at the
galactocentric radius of $0.4R_{25}$(Bresolin et al., 2016). It is also
similar to the solar metallicity of 12+log(O/H)⊙=8.66 (Asplund et al., 2005).

This archetypal barred spiral galaxy has been a showcase for multi-wavelength
studies. The Hubble Space Telescope (HST) imaged a large part of the optical
disk with both wide and narrow band filters (Blair et al., 2014). With the
proximity, these images permitted analyses of star clusters (Chandar et al.,
2010; Adamo et al., 2015; Bialopetravičius & Narbutis, 2020), showing that
high-mass clusters are less likely to form in the outer disk (Adamo et al.,
2015), and that clusters of all masses are disrupted in a timescale of a few
100 Myr in the disk (Chandar et al., 2010). The H$\alpha$ image was used to
identify very young clusters with HII regions ($\lesssim 10\,{\rm Myr}$;
Whitmore et al., 2011) and to study their interactions with the ISM (Sofue,
2018). The narrow and broad band HST images identify supernova remnants and
their progenitors (Dopita et al., 2010; Blair et al., 2014; Williams et al.,
2019). The populations of supernova remnants and HII regions, as well as their
feedback to the ISM, were also studied in radio continuum emission (Maddox et
al., 2006; Russell et al., 2020) and in X-ray (Long et al., 2014; Wang et al.,
2021). Three-dimensional optical spectroscopy covered large parts of the disk
(Blasco-Herrera et al., 2010; Poetrodjojo et al., 2019; Hernandez et al.,
2021; Della Bruna et al., 2022a, b; Long et al., 2022; Grasha et al., 2022).
These studies found a large contribution of diffuse ionized gas (DIG) to the
galaxy’s H$\alpha$ flux (Poetrodjojo et al., 2019) and its radial and
azimuthal variations: the DIG fraction is lower in the center and spiral arms
where the star formation is more active, and higher in the interarm regions
(Della Bruna et al., 2022a).

Molecular gas in M83 has been observed in CO(1-0). Early studies focused on
particular portions of the galaxy, such as the nucleus, bar, and spiral arms
(e.g., Combes et al., 1978; Wiklind et al., 1990; Handa et al., 1990; Lord &
Kenney, 1991; Kenney & Lord, 1991; Rand et al., 1999). The whole disk was
observed first with single-dish telescopes (Crosthwaite et al., 2002; Lundgren
et al., 2004b, a), which analyzed the CO $J$=2-1/1-0 line ratio between spiral
arms and interarm regions. The most recent update with the ALMA single-dish
telescopes concluded that the ratio varies systematically from $\lesssim 0.7$
in low surface density interarm regions to $\gtrsim 0.7$ in higher surface
density spiral arms when it is analyzed at a 1 kpc-scale resolution (Koda et
al., 2020).

The emission from higher CO transitions and of other molecules has also been
used to trace the gas physical conditions mainly in the central region (e.g.,
Muraoka et al., 2009a; Tan et al., 2018). These studies identified warm, dense
molecular gas in the central region (e.g., Petitpas & Wilson, 1998; Sakamoto
et al., 2004; Muraoka et al., 2009b), which could be due to gas collisions,
that accelerated the chemical evolution (Harada et al., 2019; Martín et al.,
2009). M83 was also used as a test bed of [CI] line transitions as a potential
alternative tracer of molecular regions and their masses with some success
(Jiao et al., 2021) and with caution (Miyamoto et al., 2021).

More recent CO(1-0) observations resolve molecular structures on sub-kpc
scales ($\sim 100\,{\rm pc}$; Hirota et al., 2014) and even on a cloud scale
($\sim 40\,{\rm pc}$) in a part of the disk (Hirota et al., 2018). A portion
of the eastern spiral arm, the part stretching from the bar-end toward east,
was resolved into two parallel molecular ridges. Molecular clouds along one of
the ridges show systematically higher star formation efficiency than those on
the other ridge (Hirota et al., 2018), suggesting that the star formation
trigger is related to the large-scale pattern. Egusa et al. (2018) found that
the gas velocity dispersion is enhanced in the bar, compared to the disk,
potentially spreading the gas and suppressing star formation there. The
majority of clouds without HII regions are in the interarm regions (Hirota et
al., 2018). These results, obtained in a small part of the disk, suggest that
large-scale galactic structure controls the properties of molecular clouds and
star formation. The masses of star clusters appear to be determined by those
of their parental clouds. Indeed, the mass functions of molecular clouds and
clusters appear to track each other in their upper mass cutoffs (Freeman et
al., 2017) or in their slopes (Mok et al., 2020).

The CO(1-0) data presented in this study show both large-scale galactic
structures as well as molecular clouds. It enables an unbiased census of cold
molecular gas, and mapping the cloud mass spectrum as a function of galactic
structure. The fully-resolved CO(1-0) map can potentially reveal conditions in
the immediate vicinity of star forming regions. The resolution of $\sim
40\,{\rm pc}$ is crucial for these studies, as it is the scale of a typical
molecular cloud (diameter of 40 pc), not the scale of a typical separation
between clouds ($\gtrsim 100\,{\rm pc}$; Sun et al., 2018, 2020a, 2020b;
Rosolowsky et al., 2021).

Table 1: Global Properties of M83

Parameter | Value | Ref.
---|---|---
NGC | 5236 |
Morphology | SAB(s)c | 1
R.A.(J2000) | $13^{\mathrm{h}}37^{\mathrm{m}}0\fs 57$ | 2
DEC(J2000) | $-29\arcdeg 51\arcmin 56\farcs 9$ | 2
Distance (Mpc) | 4.5 | 3
m-M (mag) | 28.25$\pm$0.15 | 3
$V_{\rm sys}$ ($\,{\rm km\,s^{-1}}$) | $511\pm 3$ | 4
$V_{\rm rot}$ ($\,{\rm km\,s^{-1}}$) | $174\pm 10$ | 4
$\Omega_{\rm b}$ ($\,{\rm km\,s^{-1}}/\,{\rm kpc}$) | $57.4\pm 2.8$ | 5
$R_{\rm CR}$ ($\,{\rm kpc}$) | 3.0 | 4,5
$M_{\rm B}$ (mag) | -20.14 | 6,3
B-V (mag) | 0.59 | 6
$D_{\rm 25}$ (arcmin) | 12.9 | 1
$D_{\rm 25}$ (kpc) | 16.9 | 1,3
Axis Ratio | 1.12 | 1
P.A. (deg) | $225\pm 1$ | 4
Inclination (deg) | $26\pm 2$ | 4
$M_{\rm star}$ (${\rm M_{\odot}}$) | $2.5\times 10^{10}$ | 7
$M_{\rm H_{2}}$ (${\rm M_{\odot}}$) | $2.6\times 10^{9}$ | 4
$M_{\rm HI}$ (${\rm M_{\odot}}$) | $7.8\times 10^{9}$ | 8
SFR (${\rm M_{\odot}}\,{\rm yr}^{-1}$) | 5.17$\pm$1.79 | 7
12+log(O/H)star | 8.78$\pm$0.07 | 9
12+log(O/H)gas | 8.73$\pm$0.27 | 9
$F_{\rm TIR}$ ($10^{-12}\,\rm erg\,s^{-1}cm^{-2}$) | 4.07$\pm$0.36 | 10

Note. — The galaxy center is obscured and is difficult to determine. We adopt
the center coordinate derived as the symmetry center of $K$-band isophotes
within the central arcmin region (Thatte et al., 2000; Díaz et al., 2006;
Sakamoto et al., 2004). The stellar and gas metallicities are the values at a
galactocentric radius of $0.4R_{25}$ ($=0.2D_{25}$). All the values are
corrected for the adopted distance when necessary. Reference: (1) de
Vaucouleurs et al. 1991, (2) Thatte et al. 2000, (3) Thim et al. 2003, (4)
This study, (5) Hirota et al. 2014, (6) Cook et al. 2014, (7) Jarrett et al.
2019, (8) Koribalski et al. 2018, (9) Bresolin et al. 2016, (10) Marble et al.
2010.

##  2 Observations

The nearby barred spiral galaxy M83 was observed with the 12m, 7m, and Total
Power (TP) arrays of ALMA in CO(1-0) under project code 2017.1.00079.S. 111The
7m and TP arrays are also called the Atacama Compact Array (ACA), or the
Morita array named after Professor Koh-ichiro Morita, a member of the Japanese
ALMA team and designer of the ACA.

Figure 1:  An optical B-V-H$\alpha$ composite image of M83 taken at the Cerro
Tololo Inter-American Observatory (CTIO) and downloaded from the NASA/IPAC
Extragalactic Database (NED). The white circle is the region observed with the
ALMA 12m and 7m arrays and encloses the whole star-forming H$\alpha$ disk of
M83. The ALMA Total Power array observed an area larger than the white circle
(see Figure 2). The yellow ellipse shows the optical edge of the disk at
$R_{25}$ (de Vaucouleurs et al., 1991).

###  2.1 12m and 7m Arrays

The whole star-forming disk of M83 was mapped with a 435-pointing mosaic using
the 12m and 7m arrays (Figures 1 and 2). Although the ALMA’s default mosaic
pattern uses different sets of pointing positions for the two arrays to
account for the different primary beam sizes, we used the same pointing
positions for the two arrays. In this way, each pointing position could be
imaged with the 12m and 7m array data separately and together, which is useful
for consistency checks.

This mosaic setup observes more positions with the 7m array than the ALMA’s
default and over-samples the area with its primary beam. However, it requires
roughly the same amount of observing time with no loss in sensitivity across
the mosaic area, as the array spends a shorter amount of time at each
pointing. The overhead associated with the over-sampling is almost entirely
from antenna slew, which is short and negligible with ALMA. Hence, the benefit
outweighs the drawback. We successfully mapped M83 in this custom manner (with
considerable support from the North American ALMA Science Center [NAASC] staff
for this setup).

The total of 435 pointing positions were split into six sub-regions (three
columns and three rows) and into 6 scheduling blocks (SBs), given the
observatory constraint of a maximum of 150 pointings per SB. Figure 2 shows
the design of the 6 SBs. The column and row arrangements overlap each other to
check consistency among SBs. For convenience, we name the SBs column 0-2 from
left to right (east to west) and row 0-2 from top to bottom (north to south).
Each of the SBs was observed multiple times, each of which is called an
execution block (EB). 20 (78) EBs were executed on the 12m (7m) array, each of
which passed the observatory quality assurance process; the breakdown is shown
in Table 2. Each EB included observations of bandpass, gain, and flux
calibrators as well as the target galaxy.

Table 2: Number of Execution Blocks (EBs)  Scheduling Block (SB) | 12m | 7m
---|---|---
column | 0 | 5 | 13
| 1 | 3 | 13
| 2 | 3 | 13
row | 0 | 3 | 13
| 1 | 3 | 13
| 2 | 3 | 13

The typical system temperatures of EBs range from 96 to 116 K, with an average
of 105 K for the 12m array, and from 91 to 138 K, with an average of 101 K,
for the 7m array. The setup of correlators (i.e., spectral windows) is listed
in Table 3. The velocity coverage and channel width are enough to cover the
full velocity width of M83 at a 0.3$\,{\rm km\,s^{-1}}$ resolution.

Table 3: ALMA Spectral Window Setups  |  | Frequency |  | Velocity
---|---|---|---|---
Array | $N_{\rm chan}$ | Band Width | Chan. Increment | Resolution |  | Band Width | Chan. Increment | Resolution
|  | (MHz) | (kHz) | (kHz) |  | ($\,{\rm km\,s^{-1}}$) | ($\,{\rm km\,s^{-1}}$) | ($\,{\rm km\,s^{-1}}$)
12m | 3840 | 234.375 | 61.035 | 122.070 |  | 609.559 | 0.15874 | 0.31748
7m | 4096 | 250.000 | 61.035 | 122.070 |  | 650.196 | 0.15874 | 0.31748
TP | 4096 | 250.000 | 61.035 | 122.070 |  | 650.196 | 0.15874 | 0.31748

###  2.2 Total Power Array

The TP observations were carried out using the On-The-Fly (OTF) mapping
technique (Mangum et al., 2007; Sawada et al., 2008). Two SBs were configured
to scan the whole galaxy at once either in the RA or DEC directions (the white
box in Figure 2), instead of the standard ALMA setup of arranging six separate
SBs corresponding to those of the 12m and 7m arrays. This minimizes scan
errors when the SBs are combined. An OFF position was integrated for 6.336 sec
every OTF scan (20.160 sec integration per raster scan). Each SB covered a
$11.7\arcmin\times 11.7\arcmin$ area, which extends beyond the molecular disk
of M83 (the white box in Figure 2). The data contain the emission-free sky on
all edges. The ACA correlator was used and the parameters are listed in Table
3.

The SBs were executed repeatedly with multiple EBs. The total of 125 EBs
obtained the Quality Assurance stage 0 (QA0) status of Pass or Semi Pass
according to the ALMA observatory records. The amplitude calibration was
performed with the standard chopper-wheel method. We rejected 5 EBs due to bad
weather (unreasonably low intensity), spurious pointing corrections (blurred
map appearance), and relatively large flux errors with respect to the other
EBs (deviations greater than a few %). We used 120 EBs after the rejections.
The average number of TP antennas participating in the observations was 3.71
out of 4. The total observing and on-source times were 110.1 h and 60.4 h,
respectively. The system temperature $T_{\rm sys}$ was 99, 105, 113 (median),
121, and 129 K at 10th, 25th, 50th, 75th, and 90th percentiles.

Figure 2:  The 435-pointing mosaic pattern of the ALMA 12m and 7m arrays. The
optical image in the background is from the Digitized Sky Survey. Each circle
represents the primary beam size of the 12m array. The pointings are split
into 6 redundant science blocks: three columns (top panel) and three rows
(bottom panel). The white box indicates the $11.7\arcmin\times 11.7\arcmin$
area mapped by the Total Power (TP) array. The 12m and 7m arrays roughly cover
a circular region with a diameter of $9.4\arcmin$ at an almost constant
sensitivity.

##  3 Data Reduction

We used the Common Astronomy Software Application (CASA: McMullin et al.,
2007; Bean et al., 2022) version 5.1 for calibration (this section), and the
Multichannel Image Reconstruction, Image Analysis, and Display (MIRIAD; Sault
et al., 1995, 1996) for imaging (Section 4).

###  3.1 12m & 7m-Arrays

The observatory performed quality assurance on each dataset and delivered a
calibration pipeline script with each data delivery. We ran them and checked
the calibrated amplitudes, phases, and fluxes of bandpass and gain
calibrators.222The ALMA observatory announced an error in amplitude
calibration in the 12m and 7m array observations, which is introduced during
the observations due to a poor calibration strategy. We neglect it because its
impact is small, and is, at most, only about 0.7% at the galactic center in
the velocity channels where the emission is the strongest. This error is
calculated as the ratio of the highest antenna temperature $T_{a}^{*}$ in the
TP cube (Section 3.2) and typical system temperature $T_{\rm sys}$. We did not
subtract continuum emission, but checked that it was not detected in the final
data cubes.

###  3.2 Total Power Array

Individual spectra were calibrated by mostly following the CASA calibration
pipeline except for the spectral baseline subtraction. We subtracted the
baselines later, from data cubes rather than from individual spectra.

The calibrated spectra were re-sampled on a grid with pixel size of
$5.62\arcsec$ using the prolate spheroidal function with a size of 6 pixels
(Schwab, 1980, 1984). The effective full-width half maximum (FWHM) beam size
after this regridding – hence smoothing – is $56.6\arcsec$. We generated
separate data cubes for the 120 EBs, calculated the flux ratios and errors of
all of their pairings, and solved for relative flux scales by inverting the
design matrix. The derived 120 scaling coefficients, whose median was set to
1, have a small scatter of only 1.3%. We applied these coefficients to co-add
all the spectra into two data cubes of the RA and DEC scans. The absolute flux
scale was monitored repeatedly by the ALMA observatory, and the consistency in
time and frequency was checked not only within our frequency band (Band 3),
but also against other bands.

Spectral baselines were subtracted with a straight line fit in these cubes.
The data reduction pipeline uses a high-order polynomial fit to individual
spectra, which sometimes causes artifacts. In practice, the baselines are
mostly flat at the observed frequency. A straight-line fit is sufficient. It
can be applied to individual spectra as done in the pipeline, or equivalently,
to data cubes after their integration, as we did in this study. The latter is
computationally much more efficient.

The RA and DEC data cubes were combined with the Emerson & Graeve (1988)
method. Figure 3 compares the noise power spectrum densities of the individual
RA and DEC scans and of their combination. Systematic noise in the scan
directions very clearly appear as excess noise power in the vertical and
horizontal directions in panels (a) and (b), respectively. This noise is
mainly due to the shared OFF integration in each OTF scan (see Heyer et al.,
1998; Jackson et al., 2006). The combination of the RA and DEC scans reduce
the noise power by a factor of about 2.

Figure 3: Power spectrum density maps of TP observations for (a) the RA scan,
(b) DEC scan, and (c) combination of the RA and DEC scans. The power spectra
of 200 emission-free channels with a channel width of 61.6 kHz are averaged.
Panels (a) and (b) include twice less integration than panel (c), and hence,
their noise power densities are scaled by 1/2 so that their color scales can
be compared directly to that of panel (c).

The antenna temperature $T_{a}^{*}$ was converted to the Jy/beam unit by
multiplying the coefficient, $C_{Jy/K}=40.7$ Jy/K, measured by the
observatory. This coefficient corresponds to a main beam efficiency of
$\eta_{\rm mb}=0.856$ via the equation

$\eta_{\rm mb}=\frac{2k_{\rm b}\Omega_{\rm
TP}}{\lambda^{2}}\frac{1}{C_{Jy/K}},$ (1)

where $k_{\rm b}$ is the Boltzmann constant, $\Omega_{\rm TP}$ is the
effective beam area, and $\lambda$ is the observing wavelength. We used a beam
size of FWHM = $56.6\arcsec$ ($\approx 2.74\times 10^{-4}\,\rm rad$) at
$\lambda=2.60\,\rm mm$, and

$\Omega_{\rm TP}=\frac{\pi\rm FWHM^{2}}{4\ln 2}.$ (2)

The final cube has a root-mean-square noise of $6.2$ mK in $T_{\rm mb}$, and
0.25 Jy/beam, in a velocity channel width of $0.159\,{\rm km\,s^{-1}}$. The
total integrated flux in the cube is $S_{\nu}dv=1.67\times 10^{4}\,{\rm
Jy\cdot km\,s^{-1}}$. This map was used to show the systematic variations of
CO $J=$2-1/1-0 ratio over the disk of M83 (Koda et al., 2020).

###  3.3 Total Power Cube to Visibilities

The TP data cube was converted to the form of interferometric data
(visibilities) using the Total Power to Visibilities package (TP2VIS; Koda et
al., 2019). This conversion permits a joint imaging (deconvolution) of the
12m, 7m, and TP data with existing imaging algorithms and software for
interferometers. By construction, the imaging algorithms work much better when
starting with a higher quality dirty image (i.e., more complete uv coverage).
The combination of 12m, 7m, and TP data before imaging ensures this higher
quality, and motivates this approach.

Two modifications were made to the method presented in Koda et al. (2019).
First, when we deconvolved the TP cube with the TP primary beam, we also took
into account the smoothing kernel introduced in re-gridding the observed data
onto the data cube grid (the first step in TP2VIS, i.e., step A in Koda et
al., 2019). That is, the TP primary beam that we used for the deconvolution is
a convolution of a Gaussian beam and the prolate spheroidal function (§3.2).

The relative weight densities among the 12m, 7m, and TP data from the
observations are not matched as a function of uv-distance (Figure 4b). Their
relative observing times are fixed by the observatory, but it causes a
shortage in sensitivity for the 7m and TP data with respect to that of the 12m
data. Thus, the weights of the 7m and TP visibilities are scaled up by factors
of 5 and 20, respectively. The factors are determined by the inspection of
Figure 4b. 333This is equivalent to down-weighting high quality data, that is,
scaling down the weights of the 12m and 7m visibilities by factors of 1/20 and
1/4 (=5/20), respectively, to match those of the TP data. Figure 4a shows the
central part of uv coverage by the 12m (blue), 7m (red), and TP visibilities
(green). Figure 4b shows their weight distributions as a function of uv
distance. The optimal relative weight distribution would show smooth
transitions from the 12m to 7m, then to TP data, which would ensure a more
optimal PSF for imaging. Figure 4c shows the distributions after the weight
scaling.

The manipulation of the imaging weights as described in the previous paragraph
may appear somewhat arbitrary. We note that the imaging of radio
interferometer data routinely involves manipulating weights. For instance,
employing robust or uniform weighting during imaging enables a user to tune
the spatial resolution and sensitivity according to the science goals of the
project. The adjustment in this study follows this precedent. Further, the
Feather task in CASA is widely used to combine TP image with 12m and/or 7m
image by implicitly adopting a weight scaling similar to ours. Feather adds
the two images without accounting for their sensitivity difference. It adjusts
the relative weights so that their combined weight densities reproduce the
ones defined by the convolution beam of the CLEANed 12m and/or 7m image (i.e.,
the restoring beam in CASA’s tclean; note that the Fourier transformation of
this beam is the weight distribution in the $uv$ space). This is equivalent to
scaling up/down the weight of the TP image. The sdintimaging task in CASA also
includes this implicit scaling.

Figure 4:  Relative weight densities of the 12m, 7m, and TP data in the uv
space for a pointing at the center of mosaic: (a) the uv distributions of 12m
(blue), 7m (red), and TP (green) data, (b) their weight distributions from the
observations, and (c) after the scaling.

##  4 Imaging

The 12m, 7m, and TP array data were jointly-deconvolved with the Steer CLEAN
algorithm (the invert and mossdi2 tasks) in the Multichannel Image
Reconstruction, Image Analysis, and Display (MIRIAD) software package (Sault
et al., 1995, 1996, this choice of software is explained in Section 4.2). The
visibility data from CASA were converted to the MIRIAD data format. For
simplicity, we adopted a single average $T_{\rm sys}$ value for each EB, as it
does not vary much within an EB in ALMA Band 3 observations. $T_{\rm sys}$ was
used for weighting in imaging. We applied the factors of 5 and 20 to scale up
the weights of the 7m and TP data as discussed in Section 3.3. The whole
435-point mosaic field was CLEANed at once. We made three data cubes with
three channel widths: 5, 2, and 1$\,{\rm km\,s^{-1}}$. We use the $5\,{\rm
km\,s^{-1}}$ cube and its parameters in the rest of the paper unless otherwise
stated. The parameters of the three data cubes are listed in Table 4.

The Briggs “robust” weight-control scheme was also employed in addition to the
$T_{\rm sys}$ weighting. We measured the sensitivity and beam size at various
robust values and adopted ${\rm robust}=-0.2$ as the optimal value for this
study. The cell size is set to $0.25\arcsec$. The root-mean-square (RMS) noise
is 3.8 mJy (96 mK) in a $5\,{\rm km\,s^{-1}}$ channel. The point spread
function (PSF) size is $2.12\arcsec\times 1.71\arcsec$ at a position angle of
$-89.0\arcdeg$.

We ran CLEAN in two steps: (1) first ran it down to a $1\sigma$ flux level
without a mask, and (2) then made a mask for large areas of contiguous pixels
at $>0.8\sigma$ significance and ran CLEAN again on the residual map from step
(1). The total flux of the residual map naturally decreased with the CLEAN
iteration. We stopped it when the average intensity within the mask hit the
floor or $\sim 0.2\sigma$. The last step was implemented because extended
emission, even at a very low level, could amount to a large total flux over
large areas. The residual map from step (1) contains about 6.6% of the total
flux of the dirty map before CLEAN, and that from step (2) contains about
1.5%.

A difference in beam area between the synthesized PSF from the uv coverage and
the gaussian PSF that is applied to clean components causes a problem in flux
measurement (Koda et al., 2019). To mitigate this problem, we matched the PSF
areas with the scheme presented by Jörsater & van Moorsel (1995). In effect,
the residual map from step (2), hence the 1.5% of the total flux, was scaled
(multiplied) by a factor based on the ratio of beam areas (0.96, Table 4) and
added to the clean component map. This step is necessary to ensure the flux
accuracy in the final data cube. The total integrated flux in the cube is
$S_{\nu}dv=1.68\times 10^{4}\,{\rm Jy\cdot km\,s^{-1}}$, which is consistent
with that of the TP data cube with a 0.6% error (Section 3.2). We refer to
this final cube as “the cube”, “the 12m+7m+TP cube”, or “the combined cube”.

Table 4: Parameters of Data Cubes

(1) | (2) | (3) |  | (4) |  | (5) |  | (6) | (7) | (8)
---|---|---|---|---|---|---|---|---|---|---
|  | Velocity |  | Beam Size |  | Noise ($1\sigma$) |  | Residual/Total | Beam Ratio | Flux Recovery
Name | $N_{\rm chan}$ | Start, End, CW |  |  $b_{\rm maj}$, $b_{\rm min}$, P.A. |  | $\delta S_{\nu}$ | $\delta T_{\rm b}$ |  |  |  |
|  | ($\,{\rm km\,s^{-1}}$) |  | ($\arcsec$, $\arcsec$, $\arcdeg$) |  | (mJy/beam) | (mK) |  | (%) |  |
CUBE5 | 100 | 250, 745, 5 |  | 2.12, 1.71, -89.0 |  | 3.8 | 96 |  | 1.5 | 0.96 | 1.01
CUBE2 | 250 | 250, 748, 2 |  | 2.12, 1.71, -89.0 |  | 5.8 | 148 |  | 2.5 | 0.96 | 1.00
CUBE1 | 310 | 340, 649, 1 |  | 2.09, 1.68, +89.9 |  | 7.7 | 195 |  | 2.7 | 0.96 | 1.00

Note. — The Briggs’s robust parameter is set to $-0.2$ for all cubes. $T_{\rm
b}{\rm\,[K]}=25.4S_{\nu}{\rm\,[Jy]}$. (1) Name of cube. (2) Number of
channels. (3) Start and end velocities, and channel width (CW). (4) Major and
minor axis diameters and position angle of the beam. (5) Root-mean-squre (RMS)
noise in a channel in units of Jy/beam and K. (6) Fraction of flux left in
residual map after the two-step CLEAN procedure. The total flux is calculated
in the dirty cube before CLEAN. (7) Areal ratio of convolution and synthesized
beams. (8) Ratio of total fluxes in 12m+7m+TP and TP cubes. The beam area
correction is applied to the residual part of the 12m+7m+TP cube. CUBE1 shows
a slightly different beam size from the other two due to the different
velocity coverage.

###  4.1 Conversion from Jy/beam to K

The data cube was converted from intensity (brightness) $I_{\nu}$, in units of
Jy/beam, to Rayleigh-Jeans brightness temperature $T_{\rm b}$ in units of K.
From the Rayleigh-Jeans equation of $I_{\nu}=2k_{\rm B}T/\lambda^{2}$ and the
relation between the intensity and flux density $S_{\nu}$ of
$I_{\nu}=S_{\nu}/\Omega_{\rm beam}$, the conversion equation is

$\displaystyle T_{\rm b}$ $\displaystyle=$
$\displaystyle\frac{\lambda^{2}}{2k_{\rm B}\Omega_{\rm beam}}S_{\nu}$ (3)
$\displaystyle=$ $\displaystyle 13.6\,{\rm K}\left(\frac{\lambda}{\rm
1mm}\right)^{2}\left(\frac{b_{\rm maj}\times b_{\rm min}}{\rm 1\arcsec\times
1\arcsec}\right)^{-1}\left(\frac{S_{\nu}}{\rm 1Jy}\right),$ (4)

where $k_{\rm B}$ is the Boltzmann constant, $\Omega_{\rm beam}$ is the beam
area in steradians, and $b_{\rm maj}$ and $b_{\rm min}$ are the FWHM sizes of
the beam along its major and minor axes in arcsec. The $S_{\nu}$ is a flux
density within a beam and can be regarded as equivalent to the $I_{\nu}$ in
the data cube.

In our observations we have $\lambda=2.60\rm\,mm$, $b_{\rm maj}=2.12\arcsec$,
and $b_{\rm min}=1.71\arcsec$, and hence, $T_{\rm b}=25.4S_{\nu}$.

###  4.2 MIRIAD Over CASA for Mosaic Imaging

We used MIRIAD for imaging. MIRIAD works better than CASA primarily because
CASA’s tclean task, as of version 6.4, does not account for spatial variations
of PSF across the mosaic, even though the uv coverage varies among the mosaic
pointing positions. Hence, CASA’s tclean uses an incorrect PSF for
deconvolution in mosaic observations. Because of this, we often ran into
runaway divergences in tclean, most likely, but not always, at the edges of
the mosaic which is covered either by the 12m or 7m array primary beam, but
not both. Some “tricks”, such as automasking or a smaller cycleniter value,
often obscure the problem, but it is unclear if they maintain accuracy since
the underlying PSF is still incorrect. Even the new task sdintimaging (Rau et
al., 2019) does not account for the spatially-variant PSF.

At this moment, we regard MIRIAD’s imaging tasks as more reliable than CASA’s
at least for a medium size or larger mosaic. The MIRIAD imaging tasks have
been used successfully in joint-deconvolution of mosaic data from
heterogeneous array interferometers and single dish telescopes (e.g., Koda et
al., 2009, 2011; Momose et al., 2010, 2013; Donovan Meyer et al., 2012, 2013;
Hirota et al., 2018; Kong et al., 2018; Sawada et al., 2018). The problem of
CASA may be less notable, or even negligible, for smaller mosaics, where all
pointing positions are observed in a sequence within a short cycle, since the
uv coverages are similar among the positions in such observations.

##  5 Maps

The reduced data is a 3-D cube with axes of right ascension, declination, and
velocity. To make 2D moment maps, we generate two mask cubes from the data
cube to include compact and extended emission components, and combined them
into a single mask cube. The first mask is defined with the regions above
$>2\sigma$ including at least one $>4\sigma$ peak in 3-D. The regions are
expanded spatially outward from their edge by one PSF major-axis diameter to
include the envelopes of the emission. This mask is designed to include
compact components detected in the original data. The second mask is made in
the same way, but with the data cube smoothed to a $5\arcsec$ resolution to
recover diffuse, extended emission around more significant emission. The
combined mask cube identifies only a few % of the pixels in the cube as
emission pixels and the rest as noise. Hence, without the mask, the moment
maps suffer significantly from the noise.

Figures 5 and 6 show the CO(1-0) integrated intensity $I_{\rm CO}$ maps (also
called the mom0 map) in color and gray in the full intensity range, and in a
limited intensity range, respectively. These three versions of the $I_{\rm
CO}$ map are presented to show diverse structures at different intensities.
Figure 6b differentiates the intensity above and below an intensity of typical
clouds in the Milky Way ($\sim 50\,\rm K\,{\rm km\,s^{-1}}$).

Figures 7 and 8 show the peak/maximum temperature $T_{\rm peak}$ map, line-of-
sight velocity field $V_{\rm los}$ (mom1) map, and the velocity dispersion
$\sigma_{v}$ (mom2) map, respectively. The mask is applied in making the
$I_{\rm CO}$, $V_{\rm los}$, $\sigma_{v}$ maps, but not for the $T_{\rm peak}$
map. We use the cube with the $5\,{\rm km\,s^{-1}}$ channel width for
generating these maps except for the $\sigma_{v}$ map, where we use the
$1\,{\rm km\,s^{-1}}$ cube so that dispersions narrower than $5\,{\rm
km\,s^{-1}}$ can be measured. The mask from the $5\,{\rm km\,s^{-1}}$ cube is
regridded and used to produce the $\sigma_{v}$ map.

The noise level $\sigma_{I_{\rm CO}}$ varies across the $I_{\rm CO}$ map
(Figure 5) depending on how many pixels along the velocity axis are included
for the velocity integration. Specifically, $\sigma_{I_{\rm CO}}=({\rm
RMS}\cdot{\Delta}V)N^{1/2}$, where RMS is the rms for channel width
${\Delta}V$ and $N$ is the number of integrated spectral channels. As a
guideline, the RMS noise of 96 mK in a $5\,{\rm km\,s^{-1}}$ channel
translates to $0.48\,\rm K\,{\rm km\,s^{-1}}$. Thereby, most emission evident
to eyes in Figure 5 is detected at a high significance. The maximum integrated
intensity is $1023\,\rm K\,{\rm km\,s^{-1}}$.

Figure 5:  The CO(1-0) integrated intensity $I_{\rm CO}$ map (the “mom0” map).
At a distance of 4.5 Mpc, an angle of $1\arcmin$ corresponds to 1.3 kpc, and
that of $1\arcsec$ corresponds to 21.8 pc. The ALMA 12m and 7m-array
observations roughly cover a $9.4\arcmin$-diameter (12.3 kpc) circular region
(see Figure 2). The beam size is $2.12\arcsec\times 1.71\arcsec$ ($46.3\,{\rm
pc}\times 37.3\,{\rm pc}$) at a position angle of $-89.0\arcdeg$. The noise
level varies spatially, as a mask is used to generate this map (as a
guideline, a pixel with only one channel integrated has an RMS noise of
$0.48\,\rm K\,{\rm km\,s^{-1}}$).  Figure 6:  The same as Figure 5, but (a) in
grayscale, and (b) with a different range of integrated intensity. These
images show faint structures clearer than Figure 5. Panel (b) shows the
regions where the integrated intensity is above and below the average within
typical molecular clouds ($I_{\rm CO}\approx 45\,\rm K\,{\rm km\,s^{-1}}$).
Figure 7:  The peak temperature $T_{\rm peak}$ maps. $T_{\rm peak}$ is the
maximum temperature/intensity value among all velocity channels at each
spatial pixel. Two panels, (a) and (b), show the same map, but with different
color scales. Panel (b) shows the low $T_{\rm peak}$ component clearer.
Figure 8:  (a) The velocity field $V_{\rm los}$ map (the “mom1” map). (b) The
velocity dispersion $\sigma_{v}$ map (the “mom2” map). Panel (a) is made from
the data cube with a $5\,{\rm km\,s^{-1}}$ channel width, while panel (b) is
from the cube with a $1\,{\rm km\,s^{-1}}$ channel width to measure
$\sigma_{v}$ smaller than $5\,{\rm km\,s^{-1}}$.

The sidelobes of a PSF often remain after imaging and are an obstacle for high
dynamic imaging, but this is not the case here. Figure 9 shows radial emission
profiles of an isolated, compact source (a) before and (b) after CLEAN. This
emission is a high velocity wing of a source at ($\alpha_{\rm J2000}$,
$\delta_{\rm J2000}$)=($13\fh 37\fm 01.0\fs$, $-29\fd 51\fm 55.4\fs$) near the
galactic center. The channels between 350 to 390 $\,{\rm km\,s^{-1}}$ show
only this emission component; these channels are integrated, and the peak
intensity is scaled to 1. The zoom-in (bottom plot) shows that there are no
large-scale, systematic sidelobes left in the CLEANed image (panel b) down to
$\sim 1/1000$.

Figure 9:  Normalized emission profiles of an isolated compact source, for an
evaluation of sidelobes on large scales: (a) for the dirty map before CLEAN,
and (b) for the CLEANed map. The peak intensity is scaled to 1. Each panel
consists of a radial profile (top) and zoom-up of $\pm 1\%$ range (bottom). A
$140\arcsec\times 140\arcsec$ cutout map around the source is the inset. The
source is a high velocity component with a size of $3.1\arcsec\times
2.1\arcsec$ (P.A.=-66$\arcdeg$) centered at ($\alpha_{\rm J2000}$,
$\delta_{\rm J2000}$)=($13\fh 37\fm 01.0\fs$, $-29\fd 51\fm 55.4\fs$). The
channels from 350 to 390 $\,{\rm km\,s^{-1}}$ are integrated. After CLEAN,
there is almost no large-scale, systematic sidelobe to the level of $\sim$0.1%
of the emission peak.

##  6 Molecular Gas Mass and Surface Density

The H2 mass and molecular gas surface density are derived from the measured CO
intensities. The total integrated CO(1-0) flux over the disk is
$S_{\nu}dv=1.67\times 10^{4}\,{\rm Jy\cdot km\,s^{-1}}$ in the TP cube (§3.2)
[$1.68\times 10^{4}\,{\rm Jy\cdot km\,s^{-1}}$ in the 12m+7m+TP cube (§4)].
This flux is translated to the total H2 gas mass of $M_{\rm H_{2}}=2.6\times
10^{9}~{}{\rm M_{\odot}}$ and total molecular gas mass of $M_{\rm
gas}=3.6\times 10^{9}~{}{\rm M_{\odot}}$ using the equations

$\displaystyle M_{\rm H_{2}}$ $\displaystyle=$ $\displaystyle 7.8\times
10^{5}~{}{\rm M_{\odot}}\left(\frac{S_{\nu}dv}{1\,\rm Jy\cdot\,{\rm
km\,s^{-1}}}\right)\left(\frac{D}{10\rm\,Mpc}\right)^{2}$ (5)
$\displaystyle\times\left(\frac{X_{\rm CO}}{2.0\times 10^{20}\,\rm
cm^{-2}\cdot[K\cdot\,{\rm km\,s^{-1}}]^{-1}}\right),$

and

$M_{\rm gas}=1.36M_{\rm H_{2}}.$ (6)

$M_{\rm gas}$ includes the masses of helium and other heavier elements that
coexist with H2. We will analyze the CO-to-H2 conversion factor with this data
in the future. For this paper, we temporarily adopt the consensus value of
$X_{\rm CO}=2.0\times 10^{20}\,\rm cm^{-2}\cdot[K\,{\rm km\,s^{-1}}]^{-1}$
(Bolatto et al., 2013) for simplicity, with caveats that this value has at
least a factor of two uncertainty (Bolatto et al., 2013), may vary with
galactic radius (Lada & Dame, 2020), and some of the measurements summarized
in Bolatto et al. (2013) are questioned (Koda et al., 2016; Scoville et al.,
2022). The optical isophotal diameter of the galaxy is $D_{25}=12.9\arcmin$
(16.9 kpc). The average molecular gas and total surface densities over the
optical disk are $\Sigma_{\rm H_{2}}^{\rm tot}\sim 12\,~{}{\rm
M_{\odot}}/\,{\rm pc}^{2}$ and $\Sigma_{\rm gas}^{\rm tot}\sim 16\,~{}{\rm
M_{\odot}}/\,{\rm pc}^{2}$.

The $M_{\rm H_{2}}$ estimated here is consistent with the previous
measurements: practically the same as $2.5\times 10^{9}~{}{\rm M_{\odot}}$ by
Crosthwaite et al. (2002) and 26% lower than $3.4\times 10^{9}~{}{\rm
M_{\odot}}$ by Lundgren et al. (2004b). For the latter, we re-calculated the
mass using the $X_{\rm CO}$ adopted in our analysis. The discrepancy is not
surprising since the telescopes, instruments, and weather for this previous
study were not as stable as the observing environment at ALMA.

The CO(1-0) integrated intensity $I_{\rm CO}$ is translated to the molecular
gas surface density as

$\displaystyle\Sigma_{\rm H_{2}}$ $\displaystyle=$ $\displaystyle 3.2\,~{}{\rm
M_{\odot}}/\,{\rm pc}^{2}\cos(i)\left(\frac{I_{\rm CO}}{\rm K\,{\rm
km\,s^{-1}}}\right)$ (7) $\displaystyle\times\left(\frac{X_{\rm CO}}{2.0\times
10^{20}\,\rm cm^{-2}\cdot[K\cdot\,{\rm km\,s^{-1}}]^{-1}}\right),$

and the total gas surface density, by including helium and other heavier
elements, as

$\Sigma_{\rm gas}=1.36\Sigma_{\rm H_{2}}.$ (8)

The $i$ is an inclination of the disk.

Under the assumption of a constant $X_{\rm CO}$, Figure 5 can be seen as the
$\Sigma_{\rm H_{2}}$ or $\Sigma_{\rm gas}$ distributions by multiplying 2.9 or
3.9, the coefficients of eqs. (7) and (8) when $i=26\arcdeg$ (derived in
Section 7). The RMS noise of $0.48\,\rm K\,{\rm km\,s^{-1}}$ in each channel
is $\Sigma_{\rm gas}\sim 1.9~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$. The maximum
integrated intensity of $1023\,\rm K\,{\rm km\,s^{-1}}$, at the inner edge of
the northern offset ridge, is $\Sigma_{\rm gas}\sim 4000~{}{\rm
M_{\odot}}/\,{\rm pc}^{2}$.

Figure 10 shows the radial profile of $I_{\rm CO}$ and $\Sigma_{\rm gas}$. The
azimuthal average is calculated in each $10\arcsec$ ($218\,{\rm pc}$) bin,
including the corrections for the inclination and position angle of the disk.
It shows an overall declining trend from the center to the disk outskirts,
with clear concentrations at the galactic center and around the bar end
($r\sim 100\arcsec$).

Figure 10:  Radial profile of integrated intensity and gas surface density.
The azimuthal average is calculated in each $10\arcsec$ bin from the galactic
center. The top and right axes are calculated assuming the distance to the
galaxy, and with the CO-to-H2 conversion factor including He and other heavy
elements. The adopted position angle and inclination are listed in Table 1.

##  7 Kinematic Parameters

Figure 8a shows that the molecular gas motion is predominantly governed by
galactic rotation, with additional local perturbations due to galactic
structures, such as the bar and spiral arms. We apply the 3D-Barolo (BBarolo)
tool (Di Teodoro & Fraternali, 2015) to the data cube and derive the kinematic
and geometric parameters of the rotating disk [i.e., systemic velocity
($V_{\rm sys}$), rotation velocity ($V_{\rm rot}$), position angle (PA), and
inclination angle ($i$)] for the rest of analysis. The derived parameters are
summarized in Table 1. A more thorough analysis of disk kinematics will be
presented elsewhere.

For stability in the fitting, we derive the parameters in three steps. Since
we need only the global disk parameters, we adopt a $10\arcsec$ radial bin to
speed up the fitting.

First, we fix the center position to the symmetry center of $K$-band isophotes
($\alpha_{\rm J2000}=13\fh 37\fm 0.57\fs$ and $\delta_{\rm J2000}=-29\fd 51\fm
56.9\fs$; Thatte et al., 2000; Díaz et al., 2006) and obtain $V_{\rm
sys}=511\pm 3\,{\rm km\,s^{-1}}$ for the radial range of $r=20$–$280\arcsec$.
We exclude the central region from the fit since it shows exceedingly high
velocity components, which are unlikely related to the galaxy’s global
dynamics.

Second, we fix ($\alpha_{\rm J2000}$, $\delta_{\rm J2000}$) and $V_{\rm sys}$,
and derive PA $=225\pm 1\arcdeg$ and $i=26\pm 2\arcdeg$ for the radial range
of $r=100$–$250\arcsec$ where the rotation curve is almost flat. The
$r\lesssim 100\arcsec$ area is excluded, because the bar ends are around
$r\sim 100\arcsec$ and non-circular motions are significant.

Third, the final rotational velocity is derived with all these parameters
fixed. The average value in $r\gtrsim 100\arcsec$ is $V_{\rm rot}=174\pm
10\,{\rm km\,s^{-1}}$, or $\Delta V=348\,{\rm km\,s^{-1}}$ in the full width.
These parameters are consistent with the ones derived in the previous CO study
(Crosthwaite et al., 2002).

We adopt the pattern speed of the bar $\Omega_{\rm b}=57.4\,{\rm
km\,s^{-1}}\,{\rm kpc}^{-1}$ from Hirota et al. (2014), which is derived from
an analysis of the geometry and kinematics of the gas in the bar offset
ridges. The corotation radius of the bar is $R_{\rm CR}=V_{\rm
rot}/\Omega_{\rm b}=3.0\,{\rm kpc}$ ($2.3\arcmin$). It is not clear if the
spiral arms have a constant pattern speed (see Section 10.3).

##  8 Molecular Gas Distribution

The integrated intensity and peak temperature maps (Figures 5-7) show numerous
local peaks across the disk, which correspond to molecular clouds or their
associations (giant molecular cloud associations - GMAs; Vogel et al., 1988).
They also show some of the common large-scale molecular structures among
barred spiral galaxies, including the concentration of molecular gas at the
galactic center, narrow offset ridges along the bar, concentrations around the
bar ends, and molecular spiral arms extending from the bar ends to the disk
outer part. Figures 5-7 also show other molecular structures, such as diffuse
extended emission around the bar and spiral arms, filamentary structures in
the interarm regions, and bifurcation or multiple branches of the molecular
spiral arms toward the disk outskirts.

In the following subsections, we will discuss the distribution and properties
of molecular gas, by separating the disk into four regions according to their
radii in the disk plane ($i=26\arcdeg$): the central region ($r<20\arcsec$),
the bar region (20-80$\arcsec$), the inner disk (80-160$\arcsec$), and the
outer disk (160-300$\arcsec$). These definitions are adapted for the purpose
of discussions in this paper, but not for rigorous classification. The regions
are indicated in Figures 11 and 12a. Figure 11 is a schematic illustration of
some molecular structures discussed in the following subsections. Figure 12
shows the probability distribution functions (PDFs) of (b) velocity dispersion
$\sigma_{v}$, (c) brightness temperature in the cube $T_{\rm b}$, (d)
integrated intensity $I_{\rm CO}$ (=$\int T_{\rm b}dv$), and (e) peak
brightness temperature $T_{\rm peak}$ (=$\max T_{\rm b}$), in these regions.

Figure 11:  Schematic illustration of molecular gas structures, drawn based on
Figures 5 and 6. The orange ellipses enclose the four regions, the central
region ($<20\arcsec$), bar (20-80$\arcsec$), inner disk (80-160$\arcsec$), and
outer disk (160-300$\arcsec$). The blue lines trace the bar and spiral arm
structures with the line thicknesses roughly indicating the significance of
emission. The blue shade shows the extended gas distributions with darker and
lighter blues for prominent and less prominent emission, respectively. The red
and brown lines are structures in the interarm regions, and the red-marked are
the ones around the transitions from the bar ends to spiral arms. This figure
does not aim to identify every structure in M83.  Figure 12:  Probability
distribution functions (PDFs). (a) Definition of circular regions around the
galactic center: the center ($r<20\arcsec$), bar (20-80$\arcsec$), inner disk
(80-160$\arcsec$), and outer disk regions (160-300$\arcsec$). The 20, 80, 160,
and 300$\arcsec$ correspond to 0.44, 1.75, 3.49, and 6.54 kpc, respectively.
The remaining panels show the PDFs for (b) velocity dispersion $\sigma$, (c)
brightness temperature $T_{\rm b}$ in the data cube, (d) integrated intensity
$I_{\rm CO}$, and (e) peak brightness temperature $T_{\rm peak}$. The PDFs are
made with the data of (b) Figure 8b, (c) the $5\,{\rm km\,s^{-1}}$ cube, (d)
Figure 5, and (e) Figure 7. The $y$ axis of each plot is a relative frequency,
proportional to the number of pixels. The $y$ scale is normalized, so that the
sum of the frequencies is one for (b), and that the frequency is one (c) at
$T_{\rm b}=2$ K, (d) $I_{\rm CO}=50\,\rm K\cdot\,{\rm km\,s^{-1}}$, and (e)
$T_{\rm peak}=2$ K.

###  8.1 Molecular Clouds in the MW Disk as Reference

The parameters of typical molecular clouds from CO surveys in the MW are a
good reference for our discussion. Table 5 lists the mass-weighted average
parameters of clouds within the solar circle from the Massachusetts-Stony
Brook Galactic Plane CO(1-0) survey (Scoville & Sanders, 1987). A typical
molecular cloud has $D\sim 40\,{\rm pc}$, $\sigma\sim 3.8\,{\rm km\,s^{-1}}$
($8.9\,{\rm km\,s^{-1}}$ in FWHM), $M_{\rm gas}\sim 4\times 10^{5}~{}{\rm
M_{\odot}}$, and $T_{\rm kin}\sim 10\,\rm K$ ($T_{\rm peak}\sim 7\,\rm K$ from
Figure 13 which is presented later). Newer measurements using the 13CO
emission provide similar cloud parameters (Roman-Duval et al., 2010; Koda et
al., 2006). We note that these mass-weighted averages are skewed toward large
and massive clouds. The clouds of $D\gtrsim 40\,{\rm pc}$ and $M_{\rm
gas}\gtrsim 4\times 10^{5}~{}{\rm M_{\odot}}$ are rare as a population, but
contain half the molecular gas mass in the MW (inside the solar circle). Most
of the clouds in the Solar neighborhood have $\lesssim 10^{5}~{}{\rm
M_{\odot}}$ (e.g., the “Taurus” cloud contains $3\times 10^{4}~{}{\rm
M_{\odot}}$; Dame et al., 1987). The gas surface densities are roughly
constant $\sim 200\,\rm~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$ among the clouds in
the inner Galactic disk, independent of their masses and sizes (Solomon et
al., 1987; Heyer & Dame, 2015), though this measure may decrease radially;
e.g., $\sim 1800\,\rm~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$ in the Galactic center
(Oka et al., 2001) and $\sim 30\,\rm~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$ in the
outer disk (Heyer et al., 2001).

Table 5: Parameters of a typical molecular cloud in the MW from Scoville & Sanders (1987) Parameter | Value
---|---
Diameter | $D$ | $40\,{\rm pc}$
Mass(H2+He) | $M_{\rm gas}$ | $4\times 10^{5}~{}{\rm M_{\odot}}$
Density | $n_{\rm H_{2}}$ | $180\,\rm cm^{-3}$
Kinetic Temperature | $T_{\rm kin}$ | $10\,\rm K$
Thermal Pressure | $P_{\rm TH}/k$ | $2000\,\rm cm^{-3}K$
Velocity Dispersion | $\sigma$ | $3.8\,\rm\,{\rm km\,s^{-1}}$

Our observations of M83 can resolve the typical clouds with $M_{\rm
gas}=4\times 10^{5}~{}{\rm M_{\odot}}$ and $D=40~{}\,{\rm pc}$ at a spatial
resolution of $40~{}\,{\rm pc}$ and $1\sigma$-sensitivities of 0.096 K in
brightness, $\sim 1.9~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$ in gas surface
density, and $\sim 4.1\times 10^{3}~{}{\rm M_{\odot}}$ in gas mass, in a
$5~{}\,{\rm km\,s^{-1}}$ channel. We can further detect, but not spatially
resolve, less massive clouds of $\sim 10^{4}~{}{\rm M_{\odot}}$ (at $3\sigma$)
if they are isolated from surrounding clouds in space and/or velocity. The
typical surface density of molecular clouds in the MW of $\Sigma_{\rm gas}\sim
200~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$ corresponds to $I_{\rm CO}\sim 45\,\rm
K\,{\rm km\,s^{-1}}$ (see Figures 5 and 6).

Most CO(1-0) emission is from optically-thick regions (see Section 9.1). Under
an assumption of thermalized, optically-thick CO(1-0) emission, $T_{\rm b}$,
$T_{\rm peak}$, and $I_{\rm CO}$ are related to the kinetic temperature
$T_{\rm kin}$ as

$T_{\rm b}=f(F(T_{\rm kin})-F(T_{\rm CMB})),$ (9)

where

$F(T)=T\left[\frac{h\nu/kT}{\exp\left(h\nu/kT\right)-1}\right].$ (10)

$T_{\rm CMB}=2.725\,\rm K$ is the temperature of the Cosmic Microwave
Background (CMB) radiation, and $h\nu/k=5.53\,\rm K$ at the CO(1-0) frequency.
$f$ is a beam filling factor, and $f=1$ when the CO emission occupies the full
area of a 40 pc beam aperture. A single, unresolved, round cloud with a
diameter of 20 pc within a beam would have $f=0.25$ and show $T_{\rm b}\sim
1.7\,\rm K$ if its kinetic temperature is $T_{\rm kin}\sim 10\,\rm K$. Figure
13 shows the relation between $T_{\rm b}$ and $T_{\rm kin}$ (eq. 9).

Figure 13:  Relation between the Rayleigh-Jeans brightness temperature $T_{\rm
b}$ and kinetic temperature $T_{\rm kin}$ at the CO(1-0) frequency. The
horizontal axis is expressed with $T_{\rm b}$ divided by the beam filling
factor $f$.

###  8.2 The Central Region: $r<20\arcsec$ ($\sim 0.44\,{\rm kpc}$)

The radial profile of $I_{\rm CO}$ (and $\Sigma_{\rm gas}$) shows a
significant concentration of molecular gas in the central region (Figure 10).
This is often observed in barred spiral galaxies (e.g., Sakamoto et al., 1999;
Sheth et al., 2005; Querejeta et al., 2021) as the bar’s elongated potential
steers gas motions toward the central region (Matsuda & Nelson, 1977; Simkin
et al., 1980; Combes et al., 1990). The central 1-kpc diameter region,
slightly larger than our definition of the central region, contains a total H2
mass of $3.9\times 10^{8}~{}{\rm M_{\odot}}$, $\sim$15% of the total in this
galaxy, with an average molecular gas surface density of $\Sigma_{\rm
H_{2}}^{\rm 1kpc}\sim 497\,~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$. [Note that for
simplicity, we assume a constant $X_{\rm CO}$ across the whole disk, while it
is suggested that the $X_{\rm CO}$ in the central region could be lower than
the adopted value (Sodroski et al., 1995; Arimoto et al., 1996; Oka et al.,
1998; Strong et al., 2004).] The $\Sigma_{\rm H_{2}}^{\rm 1kpc}$ is greater
than the average surface density of typical clouds in the inner MW disk. The
gas concentration factor, that is, the excess in surface density compared to
the disk average, is $\Sigma_{\rm H_{2}}^{\rm 1kpc}/\Sigma_{\rm H_{2}}^{\rm
tot}\sim 43$. This ratio is in the range of barred spiral galaxies studied by
Sakamoto et al. (1999) and Sheth et al. (2005) who also used a constant
$X_{\rm CO}$ value comparable to this study.

The red lines in Figure 12 show the PDFs of the central region. The excesses
in $I_{\rm CO}$, $T_{\rm b}$, and $T_{\rm peak}$ toward the high values are
evident, and are consistent with the results by Egusa et al. (2018) who
studied a smaller region of this galaxy. The high $\Sigma_{\rm H_{2}}^{\rm
1kpc}$ values likely suggest that the molecular gas occupies the entire beam.
Hence, we assume that the beam filling factor is close to $f=1$, and that the
$T_{\rm b}$ in this region is a direct measure of the $T_{\rm kin}$ of the
bulk molecular gas (see Figure 13). Figures 12c,e show that the molecular gas
there is often warmer than the temperatures in the typical Galactic clouds
($T_{\rm kin}>10\rm\,K$, i.e., $T_{\rm b}>7$ K). When averaged in a 40pc
aperture, the warmest gas shows $T_{\rm kin}\sim 19$-$21$ K ($T_{\rm b}=T_{\rm
peak}\sim 16$-$18$ K).

The velocity dispersion is also enhanced in the central region (Figure 8b). We
note that the dispersion here is primarily the component perpendicular to the
disk as the galaxy is nearly face-on. Figure 12b shows that a typical
dispersion, i.e., the peak of the PDF, is $\sigma_{v}\sim 23\,{\rm
km\,s^{-1}}$. This is much higher than those of the other regions (in the
following subsections). The excess of $I_{\rm CO}$ and $\sigma_{v}$ in the
centers of other galaxies, in comparison to their disks, is also found in
other barred spiral galaxies (Sun et al., 2020b).

M83 has a double nucleus: a $K$-band image shows a visible nucleus and
symmetry center (i.e., the bulge center; Thatte et al., 2000; Díaz et al.,
2006). The symmetry center coincides with the dynamical center (Sakamoto et
al., 2004). The highest $\sigma_{v}\sim 50\,{\rm km\,s^{-1}}$ is at the
location of the visible nucleus at (R.A., DEC)J2000=(13:37:0.95, -29:51:55.5),
which is possibly due to another gas disk rotating around this nucleus
(Sakamoto et al., 2004).

###  8.3 The Bar: $r=$20-80$\arcsec$ ($\sim 0.44$-$1.8\,{\rm kpc}$)

The most prominent features in this region are the sharp ridges of emission
along the leading sides of the stellar bar, known as “offset ridges” (dark
blue lines in Figure 11; Ishizuki et al., 1990; Kenney et al., 1992, as early
studies). These molecular ridges are coherent over a radius of $\sim 2$ kpc
and are slightly curved along the bar. Their widths are as narrow as
$\sim$100-200 pc with occasional wiggles with an amplitude of $\sim$ 100-200
pc. Their surface density is greater than the typical cloud value,
$\Sigma_{\rm gas}\gtrsim 200~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$ ($I_{\rm
CO}\gtrsim 45~{}\,\rm K\,{\rm km\,s^{-1}}$). This is also seen in the PDF as a
high $I_{\rm CO}$ tail (green in Figure 12d) compared to that of the inner
disk (blue). The offset ridges of barred spiral galaxies exhibit little star
formation (e.g., Downes et al., 1996; Momose et al., 2010; Maeda et al.,
2020), but those in M83 show some associated H$\alpha$ emission indicative of
recent star formation activity (Figure 14; see also Hirota et al., 2014).

Despite the excess in $I_{\rm CO}$ in the offset ridges, the $T_{\rm b}$ and
$T_{\rm peak}$ images show only subtle, albeit definite, differences from
those of the molecular spiral arms in the inner disk (Section 8.4). The excess
in $I_{\rm CO}$ originates almost entirely from the excess in $\sigma_{v}$
along the ridges (Figure 12b; note approximately $I_{\rm CO}\sim T_{\rm
b}\sigma_{v}$, see also Egusa et al., 2018; Sun et al., 2018). In other words,
the gas temperature is about the same between the bar and inner disk regions
even though the velocity dispersion is enhanced in the bar offset ridges. The
shape of the PDF of $\sigma_{v}$ (Figure 12b) appears to show a radial
transition from the central region (red) to the inner disk (blue). The high
$\sigma_{v}$ in the central region is also seen in the bar region (green), and
then, is reduced in the inner disk region in Figures 12b and 8b. The bar
region may carry kinematic traits of the central region. The maximum $T_{\rm
b}$ in the bar region is $\sim$10-12 K (Figure 12c,e) and occurs on the side
closer to the central region (Figure 7).

This region also shows many filamentary structures outside the offset ridges
(Figures 5 and 7; thin brown lines in the bar region in Figure 11). They are
often 1 kpc and sometimes longer in size, and some are connected to the offset
ridges almost perpendicularly on their upstream sides assuming that the disk
is rotating clockwise.

Extended, low-brightness CO emission around the offset ridges are also seen in
Figures 5 and 6 (typically $\sim 10\,\rm K\,{\rm km\,s^{-1}}$; $\sim
10^{4}~{}{\rm M_{\odot}}$ in each 40pc beam; blue shade in Figure 11). Figure
6 shows where $\Sigma_{\rm gas}$ exceeds the average cloud surface density
($\sim 45\,\rm K\,{\rm km\,s^{-1}}$) and where it is below, and hence, shows
the low-brightness emission surrounds the offset ridges. While the presence of
such extended emission was inferred in an ALMA study of the barred galaxy NGC
1300 to explain a missing flux in the interferometer data (Maeda et al.,
2020), this is the first study that such emission is clearly imaged around a
bar. The peak temperature map (Figure 7) shows that bright peaks and
ridges/filaments are spatially localized and embedded in the low-brightness
emission of much larger extents. These CO structures show only little star
formation – the extended CO emission exhibits almost no H$\alpha$ emission,
and the filaments show little H$\alpha$ emission (Figure 14). The $\sigma_{\rm
v}$ is elevated around these emission and often reaches $\sim 20$-$40\,{\rm
km\,s^{-1}}$ (Figure 8b).

Small areas around the extended, low-brightness emission show virtually no CO
emission, indicating that there is no molecular gas down to $\sim
10^{4}~{}{\rm M_{\odot}}$ within the 40 pc aperture. These areas can be more
clearly identified in Figure 8a as they appear as voids (black) in the mask
(Section 5).

###  8.4 The Inner Disk: $r=$80-160$\arcsec$ ($\sim 1.8$-$3.5\,{\rm kpc}$)

The definition of the inner disk includes the bar ends and inner sections of
the two spiral arms (dark blue and blue lines in Figure 11). The radial
profile (Figure 10) shows a notable bump around $r\sim 100\arcsec$ (2.2 kpc),
which corresponds to the transition regions from the bar ends to the beginning
of the two main spiral arms (Figure 5). The bump includes emission from the
spiral arms, as they run almost at a constant galactic radius, as well as the
emission from the gas concentrations at the bar ends. Such bar-end
concentrations are often observed in barred spiral galaxies (Sheth et al.,
2002). Compared to the long offset ridges along the bar, the gas distribution
appears more fragmented into shorter filaments and small blobs, an ensemble of
which form the concentrations at the bar ends and the spiral arm structures.
The peak brightness temperatures in the inner disk region are $\sim$8-10 K
(Figure 12c,e) and occur primarily in these concentrations (Figure 7).

The gas concentrations at the bar ends transition into two major molecular
spiral arms with a width of $\sim$$30\arcsec$ ($\sim 600\,{\rm pc}$, including
only visually-distinct bright parts of the arms). These spiral arms have a
small pitch angle and run at a near constant galactic radius (approximately
the corotation radius) in the inner disk region. In a closer look, these major
arms consist of two or more narrower filaments which run along the directions
of the major arms (see Figure 11). Each of these filaments has a typical width
of $\sim$100-200 pc. The H$\alpha$ emission is associated mainly with the
outer filaments (the convex sides), and not much with the inner filaments
(Figure 14; Hirota et al., 2018, also reported them in their study of a
smaller region).

In addition, many shorter, filamentary structures run in/out of the spiral
arms from/into the interarm regions (Figures 5 and 6; red and thin brown lines
in Figure 11). These are the structures seen as dust lanes in optical images,
dubbed “spurs” or “feathers” (Elmegreen, 1980; La Vigne et al., 2006; Chandar
et al., 2017), and are also observed in CO in other galaxies, most intensively
in M51 (Corder et al., 2008; Koda et al., 2009; Schinnerer et al., 2017).
Similar filamentary structures are also suggested in the Milky Way (Ragan et
al., 2014; Abreu-Vicente et al., 2016; Zucker et al., 2018; Veena et al.,
2021). In M83, many of the spurs are associated with HII regions (Figure 14),
which is most evident along the western spiral arm. The filamentary structures
are particularly abundant at the starting parts of both spiral arms just
outside the bar ends (red lines in Figure 11). Their symmetric occurrence at
both arms possibly suggests large-scale galactic dynamics for generating the
molecular structures.

Figure 14:  A CO(1-0) integrated intensity contour on a pseudo-color (RGB)
image from the Hubble Space Telescope. The CO contour is at 5 $\rm K\,{\rm
km\,s^{-1}}$. The HST image is made with the F657N (Hα) image for R, a
geometric mean of the F438W and F814W images (i.e., $\sqrt{\rm F438W\times
F814W}$) for G, and the F438W image for B, all of which are taken from Blair
et al. (2014).

The bright, distinct features discussed so far are surrounded by fainter,
extended CO emission (Figure 6; blue shade in Figure 11). In particular, the
emission around the western spiral arm, on its concave side, is as broad as
$\sim 50\arcsec$ ($\sim 1\,{\rm kpc}$) with a typical average surface density
of $\Sigma_{\rm gas}\gtrsim 20~{}{\rm M_{\odot}}/\,{\rm pc}^{2}$ ($I_{\rm
CO}\gtrsim 5\,\rm K\,{\rm km\,s^{-1}}$). The $\sigma_{v}$ is often as high as
$\sim 15$-$25\,{\rm km\,s^{-1}}$.

Several long spiral arm-like features are seen in the interarm regions
(outside the major spiral arms; Figure 5 and thick brown lines in Figure 11).
The most prominent is the one at the southern part, from at least around the
inner radius of the inner disk region ($r\sim 80\arcsec$ around the seven to
eight o’clock direction) to beyond the outer radius ($r\sim 160\arcsec$, five
o’clock). The full extent is difficult to trace, but it potentially stretches
inward into the bar region, and outward connected to the outer spiral arm
running along the western edge of our field of view. The length of this
feature is at least $190\arcsec$ ($3.8\,{\rm kpc}$) (only the part within the
radial range of the inner disk region) and possibly longer. The width
increases from about $15\arcsec$ (300 pc) at the inner radius to $30\arcsec$
(600 pc) at the outer radius. The dearth of H$\alpha$ emission around this
feature is remarkable, given it is a prominent molecular structure (Figure
14). Other similar filamentary features in the interarm regions are also
evident in Figure 5. Many are long ($\gtrsim 1\,{\rm kpc}$) and narrow
(100-200 pc).

A bifurcation of the western spiral arm starts in the inner disk region at
around 2 o’clock (thick blue and brown lines in Figure 11). The outer,
bifurcated branch (brown) extends into the outer disk region and is as long as
$\sim 3.5\arcmin$ (4.6 kpc). The main arm (blue) is wider and more massive
than the bifurcated arm (outside). It is notable that H$\alpha$ emission is
abundant around the bifurcated arm, but not as much around the main arm
(Figure 14). On the concave side of the main arm, there are some short
filaments running almost in parallel to, or approaching toward, the main arm
(thin brown lines in Figure 11). These filaments may collectively form a long
spiral arm-like feature in the inter arm region. This may be a counterpart of
the arm-like feature in the southern interarm region discussed in the previous
paragraph.

In addition to these prominent features, there are numerous isolated peaks
(likely individual molecular clouds) in the interarm regions (Figure 5; not
illustrated in Figure 11). These are more clearly seen in the $T_{\rm peak}$
map as blobs or dots (Figure 7).

The PDFs of $T_{\rm b}$ and $T_{\rm peak}$ in the inner disk are similar to
those in the bar (Figure 12c,e), suggesting that the physical temperature of
the gas is similar between these regions. The $I_{\rm CO}$ in the inner disk,
however, does not show the excess toward high values compared to the bar
region (Figure 12d). This is because the velocity dispersion $\sigma_{v}$ is
overall lower in the inner disk. The PDF of $\sigma_{v}$ in this region peaks
at $\sigma_{v}\sim 6.0\,{\rm km\,s^{-1}}$, substantially smaller than the one
in the central region.

###  8.5 The Outer Disk $r=$160-300$\arcsec$ ($\sim 3.5$-$6.6\,{\rm kpc}$)

Two molecular spiral arms clearly exist in the outer disk region as
continuations of those in the inner disk (Figure 5). However, they are broader
and less defined than the inner counterparts. The gas distribution appears
more flocculent in the broad spiral arms. Figure 11 illustrates them with blue
lines to show their locations, but the actual emission distribution is more
fragmented in Figure 5.

The eastern/southern arm in this region has a width of 700 pc to 1 kpc (blue
shaded area in Figure 11). It appears, in Figure 5, as a bundle of narrower
filaments or ripples, each of which is about 100-200 pc in width, running
almost parallel to each other. As a whole, these filaments form a molecular
spiral arm, which is not as confined as the molecular arms in the inner disk.
This arm shows only little H$\alpha$ emission, while only a small portion is
covered in Figure 14.

The western/northern arm runs along the outer boundary of the outer disk
region (Figure 12a and Figure 11). It is not clear if this arm is an extension
of the eastern arm in the inner disk or that of the “interarm” arm discussed
in Section 8.4 – morphologically, it also looks as if these two inner arms
merge into this outer-disk arm. Several filaments run into or out from the
main ridge of this arm but are not aligned with the arm as the filaments in
the eastern/southern arm (e.g., thinner blue lines connected this arm from its
concave side in Figure 11). The bifurcated branch of the inner western arm
discussed in Section 8.4 (thick brown line in Figure 11) also merges into this
outer arm around the 0 o’clock direction (P.A.$\approx 0\deg$). In addition to
the bifurcated branch, the feathers or spurs that emerge from the two inner
disk arms (Section 8.4) spread out into the interarm regions in the outer disk
(brown in Figure 11).

Some interarm spaces, between the main spiral arms and/or bifurcated spiral
arms, appear almost empty except several small peaks (molecular clouds). These
regions are the blank areas in Figure 8a, as they do not show significant
emission and thus are masked out (so they appear blank in the velocity field
map). In particular, the empty area around 0 o’clock is also a void in
infrared (little dust), but is prominent in UV emission (OB associations;
Jarrett et al., 2013). In fact, the HST image in Figure 14 shows almost no
H$\alpha$ emission, but shows abundant blue stars/star clusters.

The $I_{\rm CO}$, $T_{\rm b}$, and $T_{\rm peak}$ are lower in the outer disk
compared to the inner disk (Figures 5 and 7). This is quantitatively evident
in the PDFs (Figure 12). On the 40 pc scale, the molecular gas in the outer
disk is either cooler in temperature and lower in surface density than that in
the inner disk, or resides predominantly in clouds smaller than the 40 pc
resolution (or possibly a combination of these conditions). Either way, such
changes in molecular gas properties from the inner to outer disk are similar
to those in the Milky Way (Heyer & Dame, 2015). The maximum $T_{\rm b}$ values
in the outer disk region are as low as 6-8 K and occur in the wide spiral arms
(Figure 7).

The PDF of $\sigma_{v}$ in the outer disk peaks at $\sigma_{v}\sim 4.8\,{\rm
km\,s^{-1}}$, which is lower than that in the inner disk with a peak at $\sim
6.0\,{\rm km\,s^{-1}}$ (Figure 12). The difference is subtle, but
significantly detected. These typical dispersions are ubiquitous over the
inner and outer disk, and are highly supersonic (note that the sound speed in
H2 gas is $\sim 0.24\,{\rm km\,s^{-1}}$ at 10 K).

###  8.6 Masses of the Filamentary Structures

The previous subsections do not discuss the masses of the molecular structures
much, to avoid the uncertainty in $X_{\rm CO}$. Here we summarize the masses
of the filamentary structures (brown, red, and some blue lines in Figure 11)
and find that their characteristic masses are $\sim 10^{7}{\rm M_{\odot}}$.
The numbers in this section suffer from the uncertainty in $X_{\rm CO}$ (at
least a factor of 2).

In the bar region, the filamentary structures have masses of $\sim 2\times
10^{6}$ to $3\times 10^{7}{\rm M_{\odot}}$. About half of these structures
(longer ones) contain $\gtrsim 10^{7}{\rm M_{\odot}}$ (Figure 11). The inner
disk is similar. They have $\sim 3\times 10^{6}$ to $3\times 10^{7}{\rm
M_{\odot}}$. The 6 most massive have $\gtrsim 2\times 10^{7}{\rm M_{\odot}}$
and are among the red lines. The three long spiral arm-like features, or
bifurcated branches of the spiral arms, in the interarm regions (thick brown
lines) are $\sim 1\times 10^{8}{\rm M_{\odot}}$, $8\times 10^{7}{\rm
M_{\odot}}$, and $3\times 10^{7}{\rm M_{\odot}}$ for the southern, north-
western, and north-eastern features, respectively. In the outer disk, each
(blue) segment in the flocculent (fragmented) arms has a mass of $\sim 3$ \-
$7\times 10^{6}{\rm M_{\odot}}$. The two most massive segments, the most
western ones entering the western spiral arm from its concave side, have
masses of $\sim 2\times 10^{7}{\rm M_{\odot}}$ (northern one) and $\sim
8\times 10^{6}{\rm M_{\odot}}$ (southern).

##  9 Molecular Clouds and Extended CO Emission

Numerous local peaks exist across the disk of M83 (Figures 5 and 7), each of
which likely corresponds to a molecular cloud. The large molecular structures
in the spiral arms and interarm regions also appear as chains of local
peaks/clouds (Figures 5 and 7). Not all clouds are resolved, and there is
unresolved, extended emission.

A detailed analysis of the internal parameters of individual molecular clouds
will be presented in a separate paper (Hirota et al. in preparation). In
Section 9.1, we recall the basics on CO(1-0) excitation to understand the
observed emission. Most CO emission, including the unresolved emission, is
likely from molecular clouds (more accurately, from cloud-like gas
concentrations, which we call molecular clouds, independent of their internal
dynamical states, e.g., whether they are bound or not). Exciting CO emission
outside of clouds is difficult. We discuss radial and azimuthal variations of
cloud properties (Section 9.2) and a cloud-origin of the extended CO emission
(Section 9.3).

###  9.1 Collisional Excitation and Chemistry of CO

Our mass sensitivity and spatial resolution of $10^{4}~{}{\rm M_{\odot}}$
($3\sigma$) and $40\,{\rm pc}$ can detect, but cannot spatially resolve,
molecular clouds smaller than a typical one in the MW disk ($M_{\rm gas}\sim
4\times 10^{5}~{}{\rm M_{\odot}}$, $D\sim 40\,{\rm pc}$). In discussing the
full CO emission including the unresolved component, we should recall that the
collisional excitation and chemistry of CO requires volume and column
densities similar to those of molecular clouds.

In fact, the average volume density of a typical cloud ($\sim 180\,\rm
cm^{-3}$ in Table 5) is nominally not high enough for collisional excitation
of CO with H2. The critical density is an order of magnitude higher ($\sim
2,000\,\rm cm^{-3}$, Scoville & Sanders, 1987). Nevertheless, a high optical
depth within molecular cloud enables the excitation as it prevents photons
from escaping from the region. This “photon trapping” effectively reduces the
spontaneous emission rate, and as a result, decreases the effective critical
density to around the average density within clouds (Scoville & Solomon,
1974). For this reason, the bulk gas within molecular clouds can emit CO
photons even at low volume density.

Paradoxically, for molecular gas outside of molecular clouds – if it exists –
to emit significant CO emission, it has to have at least the same, or even
higher, volume density or column density than that within clouds. Such
conditions are unlikely outside of clouds except in galactic centers (Oka et
al., 1998; Sawada et al., 2001), and are unlikely to be ubiquitous across the
disk. Therefore, the great majority of the CO emission from galactic disks
should be from clouds.

In addition, there is a requirement from chemistry. CO molecules are subject
to photo-dissociation by the ambient stellar radiation field (even by the weak
field around the Sun, which is not in a major spiral arm). The presence of CO
molecules, and their emission, requires protection by self-shielding, in
addition to dust-shielding, with a sufficient column density, as well as
replenishment by efficient CO formation in a high volume density (Solomon &
Klemperer, 1972; van Dishoeck & Black, 1988). At a cloud density in a
radiation field similar to that in the solar neighborhood, the required column
density corresponds to a visual extinction of $A_{\rm V}\sim$1 mag (van
Dishoeck & Black, 1988). At a lower density, the CO formation rate decreases
($\propto{\rm density}^{2}$) faster than the photodissociation rate
($\propto{\rm density}$). Hence, a higher column density ($A_{\rm V}>$1) is
required for the self-shielding of CO. It is difficult to achieve such a
condition and to maintain CO outside molecular clouds.

Roman-Duval et al. (2016) showed that about 25% of total molecular gas mass in
the MW is detected in CO, but not in 13CO. They called it the “diffuse”
component. Goldsmith et al. (2008) performed a sensible analysis of the CO-
bright, 13CO-dark outer layer of the Taurus molecular cloud. They found a
similar mass fraction (37%) in this layer and derived the densities of
$10^{2}$-$10^{3}\,\rm cm^{-3}$. Hence, the CO-emitting “diffuse” component
across the MW is likely the cloud outer layers whose densities are as high as
the average density of molecular clouds.

Therefore, when CO(1-0) emission is detected, the majority is most likely
emitted from molecular clouds, even when individual clouds are spatially
unresolved. We note again that in this paper we use the term “molecular
clouds” for all molecular gas concentrations with the average cloud density.
We do not discuss how much of this gas is gravitationally-bound (Heyer et al.,
2001; Sawada et al., 2012; Evans et al., 2021), as it is beyond the scope of
this paper.

###  9.2 Molecular Clouds and Their Spatial Variations

The integrated intensity $I_{\rm CO}$ map (Figure 5) could portray multiple
overlapping clouds along any given line of sight, particularly in the central
region. On the other hand, the individual peaks in the $T_{\rm peak}$ map
(Figure 7) likely represent those of individual clouds. The CO(1-0) emission
is typically optically-thick, and when a cloud is spatially resolved with the
$40\,{\rm pc}$ beam (i.e., the beam filling factor is $f=1$), the observed
brightness temperature $T_{\rm b}$ (and $T_{\rm peak}$) directly reflects the
kinetic temperature $T_{\rm kin}$ of the bulk molecular gas in the cloud (see
Figure 13). The prominent molecular structures (e.g., offset ridges, spiral
arms, large interarm structures) are wider than the typical diameter of
molecular clouds, and their areas are likely filled with clouds. Hence, we
assume $f=1$ within those structures.

Figure 7 shows that $T_{\rm peak}$ decreases with galactic radius, indicating
that molecular clouds on average are the warmest in the central region and
become cooler toward the outer disk. The highest $T_{\rm peak}$ is located in
the central region (see also Figure 12e), with a maximum of $\sim
16$-$18\,{\rm K}$, corresponding to $T_{\rm kin}\sim 19$-$21\,{\rm K}$. This
is comparable to the kinetic temperatures of the clouds in the Galactic center
(Oka et al., 2001) and is twice as warm as that of typical clouds in the MW
disk (Table 5; Scoville & Sanders, 1987; Sawada et al., 2012). Beyond the
central region, the high brightness regions ($T_{\rm peak}\sim 6$-$10\,{\rm
K}$, or $T_{\rm kin}\sim 9$-$13\,{\rm K}$) are localized around the bar and
spiral arms in the inner disk (see also Figure 12e). Similarly high $T_{\rm
peak}$ values are rarely seen in the outer disk, where the highest values are
$T_{\rm peak}\sim 4$-$8\,{\rm K}$ ($T_{\rm kin}\sim 7$-$11\,{\rm K}$).

Radial variations are also seen in $\sigma_{v}$. The $\sigma_{v}$ is measured
at a cloud-scale resolution (40 pc) and likely traces the dispersion within
individual molecular clouds, except in the central and bar regions where the
line-of-sight overlap of multiple clouds may be an issue. The typical
$\sigma_{v}$ decreases from $\sim 6.0\,{\rm km\,s^{-1}}$ in the inner disk to
$\sim 4.8\,{\rm km\,s^{-1}}$ in the outer disk. Hence, the properties of
molecular clouds change with galactic radius, and $T_{\rm kin}$ and
$\sigma_{v}$ decrease radially. We note that the observed $\sigma_{v}$ is
supersonic even at the edge of the disk (the sound speed is $\sim 0.24\,{\rm
km\,s^{-1}}$ at 10 K).

Azimuthal variations in $T_{\rm peak}$ ($T_{\rm b}$) are also clear in Figure
7. In the inner disk, significant interarm structures (e.g., filaments)
exhibit the temperatures as low as $T_{\rm peak}\lesssim 3\,{\rm K}$. It may
be safe to assume $f=1$ in these structures as they maintain a high surface
density over a width of $\gtrsim 100\,{\rm pc}$ (Sections 8.4 and 8.5). If
$f=1$, $T_{\rm kin}\sim 6\,{\rm K}$ in those interarm structures, which is
lower than $T_{\rm kin}\sim 9$-$13\,{\rm K}$ in the main spiral arms in the
inner disk. Such an arm-intearm variation in molecular gas temperature has
also been reported by a CO 2-1/1-0 line ratio analysis of M83 at a lower
resolution (Koda et al., 2020).

###  9.3 Extended CO Emission

The extended CO(1-0) emission is present in and around the bar and spiral arms
(Figures 5 and 6). As discussed in Section 9.1, it is unlikely that the
majority of this CO emission comes from diffuse molecular gas outside
molecular clouds. Instead, it is likely from small, unresolved clouds.

Figure 7, as well as Figures 5 and 6, show that the low surface brightness
emission ($T_{\rm peak}\lesssim 0.5\,{\rm K}$) is ubiquitous. It is associated
with, but more extended than, the narrow molecular structures (bar, spiral
arms, interarm structures). If this extended emission consists of unresolved
clouds, the beam filling factor is $f<1$. As a thought experiment, if we
arbitrarily assume their kinematic temperature to be $T_{\rm kin}\sim 5\,{\rm
K}$ (i.e., $1\,{\rm K}$ lower than those of the significant interarm
structures, but higher than $T_{\rm CMB}$), Figure 13 gives a corresponding
value of $T_{\rm b}/f\sim 2\,{\rm K}$ (roughly consistent with the most
frequent $T_{\rm b}$ value in the Taurus cloud in the MW from Figure 4 of
Narayanan et al., 2008, especially when the large uncertainty of their main
beam efficiency for extended source is taken into account). Thus, for the
observed $T_{\rm b}=T_{\rm peak}\sim 0.5\,{\rm K}$, the filling factor in our
$40\,{\rm pc}$ beam would be $f\sim 1/4$. If the emission is from a single
cloud, its diameter is $D=20\,{\rm pc}$, and the gas mass is $M_{\rm
gas}=1.4\times 10^{4}~{}{\rm M_{\odot}}$ within a $5\,{\rm km\,s^{-1}}$
channel width. Such molecular clouds are abundant in the MW disk. While it
cannot be proven at our spatial resolution, this cloud-based explanation of
the extended CO emission seems reasonable.

##  10 Dynamical Organizations of Molecular Structures

Molecular structures in M83 are often $\gtrsim 1\,{\rm kpc}$ in length and
100-200 pc in width, having masses of $10^{7}{\rm M_{\odot}}$. They are
ubiquitous at various radii across the disk even in the interarm regions
(Sections 8.4 and 8.5). Their diversity, at first glance, appears to preclude
a unified scenario for their formation. However, with a simple thought
experiment (Sections 10.1 and 10.2), we suggest that galactic dynamics around
spiral arms play a determinant role, even for the interarm structures. The
majority of clouds should have formed in stellar spiral arms, moved out from
the parental arms without being dispersed, or remained intact after the
stellar arms were dissolved. This dynamically-driven scenario suggests that
the molecular structures move over substantial distances across the disk,
which takes time and indicates the long lifetimes of these structures and
constituent molecules once formed (Section 10.3).

In this section, we illustrate how the molecular structures can form and
evolve with basic considerations regarding their formation timescale. We lay
emphasis on the numerous filamentary structures detected in the interarm
regions. They often extend over $\gtrsim 1\,{\rm kpc}$ in length and 100-200
pc in width with a surface density enhancement of $\gtrsim 10$ in contrast to
the surrounding regions (Figures 5 and 6).

###  10.1 Disk Dynamical Timescales

For reference, we calculate two dynamical timescales using the parameters of
M83 (Table 1).

The rotation timescale of the disk is

$t_{\rm rot,disk}\sim 106\,{\rm Myr}\left(\frac{R}{3\,{\rm
kpc}}\right)\left(\frac{V_{\rm rot}}{174\,{\rm km\,s^{-1}}}\right)^{-1}.$ (11)

The rotation of the bar and spiral pattern is also important. However, a
growing number of studies support that spiral arms are dynamically-varying,
transient structures rather than the static pattern postulated by the classic
density-wave theory (D’Onghia et al., 2013; Baba et al., 2013; Dobbs & Baba,
2014; Sellwood & Masters, 2021). At least, the bar pattern must be static
since otherwise it cannot maintain its straight shape. Using a bar pattern
speed of $\Omega_{\rm p}\sim 57.4\,{\rm km\,s^{-1}}\,\,{\rm kpc}^{-1}$ (Hirota
et al., 2014), the rotation timescale of the bar is

$t_{\rm rot,pattern}\sim 107\,{\rm Myr}\left(\frac{\Omega_{\rm p}}{57.4\,{\rm
km\,s^{-1}}\,\,{\rm kpc}^{-1}}\right)^{-1}.$ (12)

###  10.2 Gas Assembly and Molecular Structures

The long lengths ($\gtrsim 1\,{\rm kpc}$) of the narrow molecular structures
(100-200 pc) already suggest the importance of large-scale galactic dynamics
in their formation. We consider a case that a long and narrow structure is
formed by a converging gas flow compressing a sheet of ambient gas in one
direction (like forming a wrinkle on a sheet). The surface densities and
widths of the ambient gas (initial condition) and of the molecular structure
(final) are $\Sigma_{\rm i}$, $l_{\rm i}$, $\Sigma_{\rm f}$, and $l_{\rm f}$,
respectively. The mass conservation gives the width of the initial region,

$l_{\rm i}=\left(\frac{\Sigma_{\rm f}}{\Sigma_{\rm i}}\right)l_{\rm f}.$ (13)

In the process, the gas has to move over a distance of $(l_{\rm i}-l_{\rm
f})$. We adopt a simplistic assumption of a coherent converging flow at a
constant velocity width of $\Delta v$ (i.e., all gas over an area of width
$l_{\rm i}$ is moving in the same directions toward a narrow structure). The
formation timescale of the structure is

$t_{\rm form}=\frac{l_{\rm i}-l_{\rm f}}{\Delta v}=\frac{l_{\rm f}}{\Delta
v}\left(\frac{\Sigma_{\rm f}}{\Sigma_{\rm i}}-1\right)\approx\frac{l_{\rm
f}}{\Delta v}\frac{\Sigma_{\rm f}}{\Sigma_{\rm i}}.$ (14)

The last approximation is for a large contrast between the formed structure
and ambient gas ($\Sigma_{\rm f}/\Sigma_{\rm i}\gg 1$).

Here we focus on how long it takes the filamentary molecular structures to
form in interarm regions, and whether such formation is possible across the
disk. In order to build up the mass of a molecular structure of width, $l_{\rm
f}\sim 100\,{\rm pc}$, and contrast, $\Sigma_{\rm f}/\Sigma_{\rm i}\sim 10$,
the initial width has to be

$l_{\rm i}\sim 1\,{\rm kpc}\left(\frac{\Sigma_{\rm f}/\Sigma_{\rm
i}}{10}\right)\left(\frac{l_{\rm f}}{100\,{\rm pc}}\right).$ (15)

The gas over a 1 kpc area must be converging coherently. Obviously, a wider
structure (e.g., $l_{\rm f}\sim 200\,{\rm pc}$) requires an even larger
initial area to sweep up the mass ($l_{\rm i}\sim 2\,{\rm kpc}$). Such a
coherent converging flow over such a large area is unlikely in the interarm
regions.

A detailed analysis of the velocity field is necessary to derive $\Delta v$,
which is beyond the scope of this paper. The $\Delta v$ is not represented by
the velocity dispersion since it does not cause a coherent convergence. The
large-scale converging flow should be due to non-circular orbital motions.

Here, we simply adopt $\Delta v=10\,{\rm km\,s^{-1}}$ as a fiducial value in
the interarm regions, surmised from a measurement in M51, a galaxy with more
prominent spiral arms. Meidt et al. (2013) derived mass-weighted azimuthally-
averaged non-circular velocities of about 5-25$\,{\rm km\,s^{-1}}$ in M51
(their Figure 2). This azimuthal average includes the spiral arms and interarm
regions, and generally, the spiral arms have larger non-circular velocities
than the interarm regions. Hence, a typical velocity in the interarm regions
is likely on the smaller side of, or less than, this range. Additionally, when
neighboring gas travels together on the same non-circular motions, their
convergence velocity $\Delta v$ should be smaller than the non-circular
velocities themselves. $\Delta v=10\,{\rm km\,s^{-1}}$ may be on a large side
within the expected range in the interarm regions.

Hence the formation of the filamentary structures in interarm regions, if they
form in-situ, requires

$t_{\rm form}\sim 98\,{\rm Myr}\left(\frac{l_{\rm f}}{100\,{\rm
pc}}\right)\left(\frac{\Delta v}{10\,{\rm
km\,s^{-1}}}\right)^{-1}\left(\frac{\Sigma_{\rm f}/\Sigma_{\rm
i}}{10}\right).$ (16)

This is as long as the rotation timescales of the disk and pattern (eqs. 11
and 12). This estimation is crude, but illustrates that their formation in the
interarm regions would take substantial times with respect to the dynamical
timescales.

Together with the required condition for a coherent converging flow ($\gtrsim
10\,{\rm km\,s^{-1}}$) over a large scale ($\gtrsim 1\,{\rm kpc}$), the
required long timescale suggests that the formation of the filamentary
structures in interarm regions is unlikely. This assessment is for general
cases and to explain the ubiquity of the structures at various radii over the
disk. Of course, some exceptional cases may still be found, but they are not
applicable in general.

Equation (16) shows that, for quick formation ($\sim 30\,{\rm Myr}$), the gas
has to move coherently at a fast speed ($\gtrsim 30\,{\rm km\,s^{-1}}$) into
the observed filamentary structures, independent of the cause of the
converging motion. In addition, these structures are massive ($\gtrsim
10^{7}{\rm M_{\odot}}$; Section 8.6), and this amount of gas has to assemble
from gas distributed over a $>1\,{\rm kpc}$ region. Observationally, it is
very rare to find such massive, large-scale coherent gas motions outside
spiral arms [unless the majority of the molecular gas is hidden in the CO-dark
phase (see Pringle et al., 2001) – however, its fraction is only $\lesssim
30\%$ within the solar radius of the MW (Pineda et al., 2013)]. This
observational constraint of a fast, coherent flow velocity of a large amount
of gas must be satisfied by any potential mechanism to be considered for the
formation of the filamentary structures.

In spiral arms, the mass naturally converges due to the large-scale
gravitational potential. With a consideration similar to the above, we could
set $\Delta v\gtrsim 30\,{\rm km\,s^{-1}}$ for spiral arms, based again on
Meidt et al. (2013). The formation timescale is shorter in the spiral arms:
$t_{\rm form}\lesssim 30\,{\rm Myr}$ for $\Delta v\gtrsim 30\,{\rm
km\,s^{-1}}$. Therefore, the filamentary structures can form in spiral arms.

Stellar feedback could also push the gas and form structures in the interarm
regions. However, the feedback-driven models so far explained only less
massive structures ($\sim 10^{6}{\rm M_{\odot}}$; Smith et al., 2020; Treß et
al., 2021; Kim et al., 2022a). Even for the small masses, these models include
spiral arm potentials and assemble the gas primarily by the potentials. The
role of feedback in the assembly process appears secondary.

###  10.3 Dynamically-Driven Evolution

The above consideration suggests that the interarm molecular structures should
have formed in stellar spiral arms. They would have traversed across the
widths of the spiral arms and flowed into the interarm regions as suggested by
the density-wave theory (Fujimoto, 1968a, b; Roberts, 1969), or remained after
the stellar spiral arms have dissolved as predicted by the swing amplification
theory (Toomre, 1981; D’Onghia et al., 2013; Baba et al., 2013). In either
case, the formation and evolution of the diverse molecular structures are
driven mainly by large-scale galactic dynamics.

The classic density-wave theory postulates steady, long-lived, stellar spiral
arms, which can accumulate coherent, large-scale molecular spiral arms. Some
parts of the molecular arms become large molecular concentrations (massive
molecular clouds; Vogel et al., 1988; Aalto et al., 1999), which can be
stretched by the shearing force along the spiral arms and by the differential
rotation after the spiral arm passages. They naturally extend to
spurs/feathers in the interarm regions (Koda et al. 2009; see also Corder et
al. 2008; Schinnerer et al. 2017). In the inner disk of M83, the molecular gas
spiral arms appear relatively focused as narrow ridges (about 600 pc widths).
The spurs/feathers are resolved as chains of molecular clouds (Koda et al.,
2009), and often rooted to, and extending out of, the main spiral arms.
Notably, they are found mostly in the inner disk, predominantly on the convex
sides of the spiral arms (Section 8.4) as expected in the density-wave type of
gas flow.

The swing amplification model predicts transient, short-lived stellar spiral
arms, where the gas also accumulates, dissipates, and becomes denser to form
molecular structures (Baba, 2015; Baba et al., 2015). The stellar spiral
structures are temporary density enhancements, but self-perpetuate themselves
by forming subsequent spiral arms nearby (D’Onghia et al., 2013). The
temporary nature results in less coherent gas structures (Baba, 2015; Baba et
al., 2015), which can be left after the stellar spiral arms dissolve. In the
outer disk of M83, the molecular spiral arms continue from the inner disk, but
are broader, less spatially coherent, and appear flocculent.

Baba (2015) suggested a gradual radial transition in a barred spiral galaxy,
from the density-wave picture in the inner disk to the swing amplification
picture in the outer disk. In their model, even the inner spiral arms are
triggered by the swing amplification, but live longer due to the gravitational
influence of the static bar potential, and behave like steady density waves
for some period of time (i.e., a good fraction of the disk rotation timescale
- eq. 11). The observed transition in M83 from the coherent to less-coherent
molecular spiral arms from the inner to outer disk may fit to this transition
picture.

The bar region also shows filamentary molecular structures. Some of these
features run into the offset ridges from their upstream sides, assuming that
the disk is rotating in the clockwise direction (Figure 5). They could also be
sheared-off remnants that traveled from the previous offset ridge. Given a
relatively short crossing timescale of $\sim 20$-$80\,{\rm Myr}$ (Hirota et
al., 2014, their Figure 17), their remnants can reach the next offset ridge
before they are completely broken apart. Such a formation mechanism of
stretched filamentary structures has been discussed with simple orbit models
in bar potentials (Koda & Sofue, 2006; Hirota et al., 2014).

The discussions here are based only on a basic estimation of their formation
timescale and on their morphologies. It is not surprising that some exceptions
exist. For example, isolated molecular clouds in the interarm regions, which
appear as dots in Figures 5-7, could have formed in local density
fluctuations, rather than a result of galactic dynamics.

###  10.4 Implications on Cloud Lifetimes

The discussion above suggests that the large molecular structures survive
through the dynamical processes in the galaxy and have lifetimes of an order
of $\sim 100$ Myr. Since the CO(1-0) emission is primarily from molecular
clouds (see Section 9.1), it implicates that the lifetimes of the embedded
molecules and molecular clouds are also as long.

Here, as a definition of molecular cloud lifetimes, we adopt the full duration
that the gas is molecular in molecular clouds, and hence, is in a prerequisite
condition for potential star formation. It has been suggested that molecular
clouds evolve through coagulation and fragmentation (Scoville & Hersh, 1979;
Vogel et al., 1988; Koda et al., 2009; Dobbs & Pringle, 2013), and that some
parts within clouds (e.g., dense cores) could be dispersed by stellar feedback
(e.g., Elmegreen, 2007). Those processes might be acting within the observed
molecular structures. Some authors defined cloud lifetimes as a short branch
period between a coagulation and fragmentation (e. g., Kruijssen et al., 2019;
Chevance et al., 2020; Kim et al., 2022b), which is a correct definition in
its own way. However, most cloud studies stem from interests in star
formation, and it seems more relevant to adopt the full duration that the gas
can potentially form stars as cloud lifetimes.

##  11 Summary and Conclusions

We present an ALMA imaging of molecular gas across the full star-forming disk
of the barred spiral galaxy M83 in CO(1-0). The joint deconvolution of the
data from the ALMA 12m, 7m, and Total Power arrays was performed with the
MIRIAD and TP2VIS packages. The mass sensitivity of $10^{4}~{}{\rm M_{\odot}}$
is sufficient to detect the most abundant population of molecular clouds with
masses $\lesssim 10^{5}~{}{\rm M_{\odot}}$. The spatial resolution of 40 pc is
the typical physical diameter of clouds in the MW ($D\sim 40~{}\,{\rm pc}$).
Therefore, this case study of one galaxy is complementary to the large survey
of nearby galaxies in the excited CO(2-1) transition (Leroy et al., 2021b, a),
which mainly detects a population of massive molecular clouds of $\gtrsim
10^{5}~{}{\rm M_{\odot}}$ (Rosolowsky et al., 2021) at a spatial resolution of
45-120 pc, high enough to separate those clouds, but not to resolve the
typical cloud diameter (Sun et al., 2018, 2020b).

The molecular gas distribution shows coherent large-scale structures in the
inner part, including the central gas concentration, offset ridges along the
bar, and prominent molecular spiral arms in the inner disk. In the outer disk,
the molecular spiral arms are still present, but appear less coherent and
flocculent. Massive filamentary gas concentrations are present both in the
spiral arms and interarm regions in the inner and outer disks, as well as in
the bar region. All of these structures embed molecular clouds and appear as
chains of clouds. Many of the interarm structures host little or no star
formation traced by H$\alpha$ emission. Unresolved, extended CO emission is
also detected around these structures.

The data show the radial and azimuthal variations in molecular gas properties,
from the galactic center, bar, and inner to outer disks, and between the
spiral arms and interarm regions. The local properties, such as brightness
temperature, velocity dispersion, integrated intensity, and hence, surface gas
density, decrease outward. They are higher in the spiral arms and lower in the
interarm regions. Given that the spatial resolution is comparable to the
clouds’ typical diameters, these radial and azimuthal variations can be
attributed to the variations of clouds’ internal properties. The majority of
the detected CO(1-0) emission, including the unresolved CO emission, is most
likely from molecular clouds (cloud-like gas concentrations).

We describe a scenario in which the ubiquitous large molecular structures,
especially the ones in the interarm regions, can form and evolve. It would
take too long to assemble their huge masses without coherent converging gas
flows over very large areas. Such flows are expected around spiral arms, but
not in the interarm regions. The interarm structures are therefore unlikely to
form in-situ. Instead, we suggest that they assemble within the stellar spiral
potential, and are either expelled into the interarm regions from the spiral
arms, or are left behind after the stellar arms dissolve. This indicates that
the molecular structures, and embedded molecular clouds, survive through these
dynamical processes and have lifetimes of the same order as the rotation
timescale of the disk ($\gtrsim 100$ Myr). The overall distribution of the
molecular structures appears to be consistent with the gradual radial
transition from the spiral density-wave picture in the inner disk to the swing
amplification picture in the outer disk.

This picture of dynamically-driven molecular gas evolution is suggested to
explain the majority of the molecular structures and clouds observed in M83.
The discussions are based only on a basic estimation of their formation
timescale and on their morphologies. It is not surprising that some exceptions
exist. We expect that subsequent studies and future observations will confirm
or modify this picture.

We thank the anonymous referee for useful comments. We thank all staff members
at the Joint ALMA Observatory (JAO) and North American ALMA Science Center
(NAASC), who helped us customize the observation strategy and made the data
reduction possible. JK also thanks Peter Teuben for his help with MIRIAD and
TP2VIS. This paper makes use of the following ALMA data:
ADS/JAO.ALMA#2017.1.00079.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 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 Extragalactic
Database (NED), which is operated by the Jet Propulsion Laboratory, California
Institute of Technology, under contract with the National Aeronautics and
Space Administration. JK acknowledges support from NSF through grants
AST-1812847 and AST-2006600, and from NAASC through the ALMA Development Study
program to develop TP2VIS. FE is supported by JSPS KAKENHI grant No. 17K14259.
AGdP has been partly supported by grant RTI2018-096188-B-I00 funded by
MCIN/AEI/10.13039/501100011033. LCH was supported by the National Science
Foundation of China (11721303, 11991052, 12011540375, 12233001) and the China
Manned Space Project (CMS-CSST-2021-A04, CMS-CSST-2021-A06).

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