# A Multiwavelength Study of the Sgr B Region: Contiguous Cloud-Cloud
Collisions Triggering Widespread Star Formation Events?

Rei Enokiya Department of Physics, Nagoya University, Chikusa-ku, Nagoya,
Aichi 464-8601, Japan  Department of Physics, Faculty of Science and
Technology, Keio University, 3-14-1 Hiyoshi, Kohoku-ku, Yokohama, Kanagawa
223-8522, Japan  Yasuo Fukui  Department of Physics, Nagoya University,
Chikusa-ku, Nagoya, Aichi 464-8601, Japan

(Accepted ApJ)

###### Abstract

The Sgr B region, including Sgr B1 and Sgr B2, is one of the most active star-
forming regions in the Galaxy. Hasegawa et al. (1994) originally proposed that
Sgr B2 was formed by a cloud-cloud collision (CCC) between two clouds with
velocities of $\sim$45 km s-1 and $\sim$75 km s-1. However, some recent
observational studies conflict with this scenario. We have re-analyzed this
region, by using recent, fully sampled, dense-gas data and by employing a
recently developed CCC identification methodology, with which we have
successfully identified more than 50 CCCs and compared them at various
wavelengths. We found two velocity components that are widely spread across
this region and that show clear signatures of a CCC, each with a mass of
$\sim$106 M⊙. Based on these observational results, we suggest an alternative
scenario, in which contiguous collisions between two velocity features with a
relative velocity of $\sim$20 km s-1 created both Sgr B1 and Sgr B2. The
physical parameters, such as the column density and the relative velocity of
the colliding clouds, satisfy a relation that has been found to apply to the
most massive Galactic CCCs, meaning that the triggering of high-mass star
formation in the Galaxy and starbursts in external galaxies can be understood
as being due to the same physical CCC process.

Interstellar medium(847); Molecular clouds (1072); CO line emission (262);
Galactic center (565); Star formation (1569)

##  1 Introduction

###  1.1 The Center of the Galaxy

Despite its small volume, the inner 1 kpc of the Galaxy contains several to
10$\%$ of molecular gas, and thus it is a prominent mass reservoir (Morris &
Serabyn, 1996). As expected from the large amount of gas, the strength of the
frozen-in magnetic fields is very high in the Galactic Center (GC). It has
been estimated to be at least 30 $\mu$G (Crocker et al., 2010) and this leads
to very large Alfvén speeds of a few tens to a hundred km s-1. Thus, the
nonnegligible field can accelerate gas, and this makes the situation more
complicated in the GC (Fukui et al., 2006).

Molecular-line observations have revealed that the molecular gas in the GC
exhibits a very large velocity width and a parallelogram shape in the $l$–$v$
diagram, which cannot be explained by Galactic rotation (e.g., Oka et al.,
1998; Tsuboi et al., 1999). The acceleration mechanism of the gas is still
unclear, but so far two possible accelerators have been suggested.

One is the bar potential, originally proposed by Binney et al. (1991). They
suggested that the framework of the parallelogram in the $l$–$v$ diagram is
explained by the innermost orbits of the so-called $x_{1}$ orbits. However,
there has been no firm observational evidence of the small bar, which is
required to fit the size of the parallelogram to the observed one.
Furthermore, many recent theoretical and observational works have revealed
that the bar is much larger and rotates at a much slower pattern speed than
assumed in Binney et al. (1991) (see e.g., Wegg & Gerhard, 2013; Sormani et
al., 2015; Portail et al., 2017; Sanders et al., 2019).

In accordance with those findings, Fux (1999) performed simulations using a
composite $N$-body and hydro code, and exhibited asymmetric and non-stationary
gas flows that substantially reproduce the observed $l$–$v$ diagrams of Hi and
CO. Recent hydrodynamic simulations by Rodríguez-Fernández (2011) and Sormani
et al. (2018) have successfully explained the lateral sides of the
parallelogram as dust lanes or gas that has been collisionally accelerated by
dust lanes. On the other hand, the top and bottom sides of the parallelogram
are still unclear. Although Sormani et al. (2018) have interpreted them as
brushed gas between the dust lanes and the central molecular zone (CMZ), the
boundaries of the top and bottom sides in the simulations still remain
obscured and are not so clear compared to the CO emissions seen in Oka et al.
(1998).

The other acceleration mechanism involves magnetic instabilities, such as the
magnetorotational instability and Parker instability originally proposed by
Fukui et al. (2006). Some magnetohydrodynamic (MHD) simulations have revealed
that this mechanism can also accelerate gas and can reproduce a parallelogram-
like distribution as a transient feature (e.g., Suzuki et al., 2015), although
the shape is blobby and not clear. This model also has observational
difficulty. Obtaining a direct detection of the magnetic field embedded in the
dense molecular gas with a large line width is fairly difficult, so there have
been few works that have tackle the situation through a new method, using far-
infrared (FIR) polarimetry and velocity gradients of gas (Hu et al., 2021).

Similar to the acceleration mechanism, many models for the dynamical structure
of the molecular condensation within the central 200 pc—hereafter, the CMZ
(Morris & Serabyn, 1996)—have been proposed (e.g., an expanding shell—Sofue
1995; a twisted ring—Molinari et al. 2011; and an open orbit—Kruijssen et al.
2015). Regardless of which model is ultimately proven to be correct, revealing
the star formation mechanism in the dense molecular cloud in such an extreme
magnetized, highly pressured environment is important for our understanding of
the Galaxy evolution.

Although the central mass reservoir often acts as a star formation engine,
star formation in the GC is inactive compared to the theoretical expectation,
and possibly suppressed (e.g., Longmore et al., 2013). In order to unveil this
star formation mystery, a multiwavelength study—including both dense- and
diffuse-gas tracers—will be important, because the star formation timescale is
shortened in such extreme condition by external pressures (Kauffmann, 2016),
and hence the dense gas no longer holds information regarding the triggering
mechanism. Thanks to the improvement in heterodyne receivers over the last two
decades, plenty of molecular line surveys have been carried out (e.g., Jones
et al., 2012). However, most works focus only on the main streams of the CMZ,
with dense-gas tracers, and do not include the diffuse gas traced by CO. We
therefore focus on the most active star forming region in the CMZ, Sgr B, and
conduct a multiwavelength study with the latest archival data sets, which
cover large area and density range, integrating them with knowledge obtained
from previous studies.

###  1.2 The Sgr B region

The Sgr B region, consisting of Sgr B1 and B2, is one of the brightest radio
sources, and is located a projected distance of $\sim$100 pc away from the GC.
The Sgr B region is also known to be a strong SiO source in the Galaxy
(Martín-Pintado et al., 1997). Mehringer et al. (1992) carried out
observations of the radio-continuum and recombination lines toward Sgr B1
using the Very Large Array (VLA), and they found extended emission in Sgr B1,
unlike Sgr B2. They discussed the possibility that the emission may have been
caused by coalesced, evolved Hii regions, and from the proximities of Sgr B1
and Sgr B2 in terms of their projected distances and revealed ionized gas
velocities, they suggested that Sgr B1 and Sgr B2 are a unified structure.
Mehringer et al. (1993) investigated OH and H2O masers toward Sgr B1 using the
VLA, and discovered an OH maser and ten H2O maser spots. These observations
provide evidence for ongoing star formation in Sgr B1, although the majority
of the radio-continuum emission—or, in other words, the star-forming
activity—in this region is dominated by evolved Hii regions. Although
estimating a quantitative value for the star-forming activity in the GC region
is very difficult, due to the heavy extinction at optical wavelengths, Simpson
et al. (2018) found that at least eight separate subregions were highly
excited by high-mass stars. From their analyses of the spectral energy
distributions obtained from [Oiii] 52 and 88 $\mu$m data, taken with the
Stratospheric Observatory for Infrared Astronomy, they concluded that these
high-mass stars were equivalent to late O-type stars. They also estimated the
ages of those late O-type stars to be at least a few Myr. From analyses of
radio-continuum emissions obtained with the Wilkinson Microwave Anisotropy
Probe, Barnes et al. (2017) also estimated the total stellar mass embedded
within Sgr B1 to be $\sim$8000 M⊙.

Sgr B2 is young, and it is the most active star-forming region in the Galaxy.
It contains numerous O-type stars, Hii regions, ultracompact and compact Hii
regions, H2O masers (e.g., Yusef-Zadeh et al., 2004; McGrath et al., 2004),
and also a wealth of chemical species, such as organic molecules. The IR
luminosity is incredibly high, $\sim$107 L⊙ (Goldsmith et al., 1987), and its
FIR spectrum and [Cii] 157.7 $\mu$m line are comparable to those in FIR-bright
galaxies (Vastel et al., 2002). Almost all of the star-forming activity is
concentrated within three regions, with diameters of $\sim$1 pc, called Sgr
B2(N), Sgr B2(M), and Sgr B2(S), which are aligned from north to south (Benson
& Johnston, 1984). Sgr B2(M) is a notable starburst region that contains 49
ultracompact Hii regions (Gaume et al., 1995). Recently, Ginsburg et al.
(2018) carried out 3 mm radio-continuum observations with ALMA, and discovered
271 compact sources, which may consist of hypercompact Hii regions and massive
young stellar objects, within a 5 pc radius around Sgr B2(M).

Sgr B2 is associated with a giant molecular cloud containing 6 $\times$ 106 M⊙
(Lis & Goldsmith, 1989). Sofue (1990) discovered giant molecular complexes
associated with Sgr B1 and B2, namely, the Sgr B1 molecular shell and the Sgr
B2 molecular complex. He suggested that star formation in Sgr B2 was triggered
by a compression induced by the expansion of the Sgr B1 molecular shell,
driven by feedback from Sgr B1. On the other hand, Hasegawa et al. (1994)
proposed a cloud-cloud collision (hereafter, CCC) scenario for the Sgr B2
region, based on their 13CO($J$=1–0) observations with the Nobeyama 45 m
telescope. They discovered three clouds in the direction of this region, which
they named the Shell, at $V_{LSR}$ = 20–40 km s-1; the Hole, at $V_{LSR}$ =
40–50 km s-1; and the Clump, at $V_{LSR}$ = 70–80 km s-1 (see also Figure 1).
They interpreted the Clump, Hole, and Shell as clouds that have been
compressed, hollowed, and shocked by the CCC, respectively, and they concluded
that massive star formation in the Sgr B2 complex was triggered by the
collision. Sato et al. (2000) and Hasegawa et al. (2008) lent further support
to this model, by comparing the spatial and velocity distributions of their
C18O($J$=1–0), C18O($J$=3–2), CS($J$=1–0), and CS($J$=7–6) data sets with
those of the masers. However, these models are based on under-sampled
datasets, so the effective spatial resolution is low. Also, this scenario
conflicts with recent observational results at other wavelengths. By analyzing
HNCO data, Henshaw et al. (2016) found that there is only one cloud toward Sgr
B2, and they suggested that the cloud identification by Hasegawa et al. (1994)
was suspicious. Based on a dynamical model of the CMZ, Kruijssen et al. (2019)
suggested that the complex kinematic structure of Sgr B2 was naturally
explained by the superimposition of fragmentation and orbital motion of a
cloud in the line of sight (LOS).

Although Sgr B2 is the most active star-forming region in the Galaxy, the
other active star-forming region that physically connects with it—i.e., Sgr
B1—has not been explained together with Sgr B2 in either scenarios.
Furthermore, the CCC scenario has been gradually taken into account in recent
dynamical models of the CMZ. Sormani et al. (2020) suggested that CCCs were
naturally present in their hydrodynamic simulations, thus it is important to
examine whether the observational footprints of the CCCs are consistent with
those simulations and dynamical models. Therefore, reexamination of the CCC
scenario for Sgr B2 based on existing dynamical models is required, and this
scenario also needs to be applicable mot only to Sgr B2 but also to Sgr B1.

###  1.3 CCC and High-mass Star Formation

It is well-known that a starburst in a merger is triggered by a collision
between two giant molecular clouds belonging to the two galaxies (see, Young
et al., 1986). A CCC is a mechanism that collects gas and magnetic fields at
super sonic velocity. Numerical simulations indicate that a CCC produces a
compressed layer of gas that generates massive, gravitationally unstable,
molecular clumps within it. Owing to the enhancement of the magnetic field
strength by shock compression, the clumps have effective Jeans masses large
enough to develop into high-mass stars (Habe & Ohta, 1992; Inoue & Fukui,
2013; Inoue et al., 2018). However, it is still unknown how many new stars are
triggered by a CCC. In other words, the controlling parameters, which
determine the scale of the star-forming activity initiated by a CCC, are still
unknown, although they must be related to the initial physical conditions of
the collision.

Thanks to recent Galactic plane surveys with 1 pc to subpc resolution—such as
those done by the Mopra and Nobeyama 45 m telescopes in CO (Braiding et al.,
2018; Umemoto et al., 2017)—the mechanism operating in CCCs have gradually
been unveiled. These datasets exhibit many candidate CCCs in the Galaxy; in
total, more than 50 Hii regions and star clusters are possible sites of CCCs
(e.g., Nishimura et al., 2018; Fukui et al., 2018c, b; Ohama et al., 2018;
Enokiya et al., 2018; Kohno et al., 2018; Torii et al., 2018; Sano et al.,
2018; Hayashi et al., 2018). Also, nearby extragalactic giant Hii regions,
such as NGC 604 in M 33 and R136 in the Large Magellanic Cloud, have been
suggested to be the results of large-scale collisions between two giant Hi
clouds (Tachihara et al., 2018; Fukui et al., 2017). These results suggest
that a CCC may be the major mechanism responsible for triggering high-mass
star formation (Fukui et al., 2021).

In order to link our understanding of CCCs in starburst galaxies with those in
the Galactic Hii regions, the nearest center of the galaxy—the GC of the Milky
Way—is the most important target. The GC has the highest volume density of
molecular clouds in the Galaxy, and hence, it is likely to have the highest
frequency of CCCs as well (Enokiya & Fukui, 2021). The physical parameters of
molecular clouds in the GC, such as their densities and relative velocities,
are different from those in the Galactic disk. The GC is thus a suitable
region for investigating these control parameters. As a first step for
investigating CCCs in the GC, Enokiya et al. (2021a) applied a CCC
identification methodology developed by Fukui et al. (2021) to the common
footpoint of magnetic loops 1 and 2, and they found two pairs of colliding
clouds. Thus, they not only demonstrated the validity of this methodology in
the GC, but also determined a unique property only seen in CCCs in the GC:
they detected SiO emissions caused by a larger relative velocity of $\sim$a
few tens of km s-1.

###  1.4 The Aims of This Paper

The CCC scenario in Sgr B2, originally proposed by Hasegawa et al. (1994) and
followed by Sato et al. (2000), has attracted much attention. However, there
are still some uncertainties, as below. This scenario:

  1. 1.

selected three velocity clouds by eye, but this is arbitrary;

  2. 2.

was only based on the distribution of molecular gas, and did not consider
other wavelengths;

  3. 3.

used undersampled data sets;

  4. 4.

neglected Sgr B1; and

  5. 5.

did not consider the dynamical structure of the CMZ, leaving the possibility
of clouds colliding at the position of Sgr B2 uncertain.

In the present paper, we utilize the CCC identification method recently
developed by Fukui et al. (2021) and Enokiya et al. (2021a), employing recent,
fully sampled data sets with large enough coverage to include Sgr B1 and large
enough density dynamic range to investigate whether the CCC model is still
applicable. By comparing the gas distribution and recent high-resolution IR
and radio-continuum data, we have constructed a multiwavelength view of the
region. We finally discuss the possibility of the CCC event and resulting star
formation, together with the dynamical structures of the CMZ.

Sections 2 and 3 provide information about the methodology and data sets used
in this paper. Our main results are described in Section 4. Section 5 is
divided into three sub-Sections, in which we discuss comparisons with
numerical simulations (5.1), a possible scenario based on a CCC (5.2), and
comparisons with other CCCs in the Galaxy (5.3). We summaries the present
study in Section 6.

##  2 Methodology

###  2.1 Signatures of a CCC

In this Section, we first introduce the methodology developed by Fukui et al.
(2018a, 2021) and Enokiya et al. (2021a, b) to identify a CCC. This
methodology has so far successfully identified more than 50 CCC candidates.
Based on synthetic observations of numerical simulations (Takahira et al.,
2014; Haworth et al., 2015) and on observational data of the Galactic Hii
regions, Fukui et al. (2021) proposed the following signatures of a CCC;

  1. 1.

Complementary distributions between two colliding clouds: a larger cloud is
hollowed out by a smaller cloud during a collision and forms the complementary
spatial signature. This signature is sometimes accompanied by a spatial
displacement, depending on the inclination angle of the relative motion to the
LOS.

  2. 2.

A V-shaped structure or bridge feature in the position-velocity diagram: In
the early stage of a CCC, the contact surface between the two clouds forms
bridge features by exchanging momenta. This signature develops into a V-shaped
structure when the larger cloud is hollowed out.

  3. 3.

Positional coincidence between the compressed layer and the high-mass star-
forming region.

However, not all of these signatures are always observable, because the
signatures are weakened by dispersion of the natal clouds caused by the CCC
and by ionization and projection effects (Fukui et al., 2021). In addition,
Enokiya et al. (2021a) found that CCCs in the GC are accompanied by SiO
emissions, due to the larger relative velocity achieved only in the GC. We
utilize these observational signatures to identify a CCC.

###  2.2 The Cold Gas Tracer

A gas density tracer is quite important for utilizing this methodology. Sgr B2
has an extremely high column density, 4 $\times$ 1025 cm-2 (Molinari et al.,
2011), and emissions from this region often show self-absorption. Thus, the
region is usually traced by rare molecules such as HNCO (e.g., Henshaw et al.,
2016). On the other hand, diffuse low-density gas, which is the dominant mass
reservoir, and mainly traced by CO, is also important for this work, because a
CCC is a mechanism for compressing the gas, and both cause and effect—in other
words, diffuse and dense gases—are required to identify a CCC. Fukui et al.
(2021) suggested that C18O is the best tracer for the condition of Sgr B2.

We carefully checked various lines obtained with previous surveys, not only
toward the Sgr B region, but also the entire CMZ, and found almost all lines,
such as 12CO, 13CO, SiO($J$=1–0), H13CO+($J$=1–0), CS($J$=1–0), HCN($J$=4–3),
H13CN($J$=1–0), HNC($J$=1–0), etc., were optically thick toward Sgr B2 and Sgr
A. On the other hand, a few lines, i. e., C18O, SiO($J$=5–4), HC3N($J$=10–9),
and HNCO(4–3) did not show significant self-absorption even toward Sgr B2.
Within these optically thin lines, SiO and HC3N($J$=10–9) are a shock tracer
and warm gas tracer, respectively, and do not show any diffuse emissions
traced by CO. Thus, they do not trace density of cold molecular gas. By
comparing a high sensitivity H2 column density map derived from $Herschel$
data (Molinari et al., 2011), we found C18O($J$=2–1) to be the best gas
tracer, with the largest density dynamic range of density. HNCO is a good gas
tracer, also including the diffuse component but the distribution of Sgr B2 is
apparently different from that of the H2 column density. This is possibly
caused by the special chemical environment in Sgr B2(N), where rare molecules
are very abundant (e.g., Belloche et al., 2013). Also, the sensitivity of
archival C18O data was much better than that of the HNCO data. In this paper,
therefore, we mainly use C18O($J$=2–1) as the molecular-gas tracer for
identifying a CCC in Sgr B2.

###  2.3 Cloud Identification

One of the signatures of a CCC, i. e., a V-shaped structure or bridge feature
in the position-velocity diagram, as introduced above, can also be seen in the
velocity gradient of a cloud. Thus, identifying two independent clouds is
important factor for the CCC methodology; whereas two clouds are often merged
by the collision and are observed as a single cloud, at least partially. We
utilized a cloud separation method, using moment maps introduced by Enokiya et
al. (2021b). This method can clearly separate two merging clouds in terms of
velocity, and it has been successfully applied to GC clouds (Enokiya et al.,
2021a). The detailed justification of this method is described in the Appendix
in Enokiya et al. (2021b).

##  3 Data Sets

###  3.1 Molecular Lines

As the main tracer of molecular gas, we use the public C18O($J$=2–1) data
obtained by Ginsburg et al. (2016), taken in the on-the-fly mode with the
Atacama Pathfinder Experiment (APEX). The spectra in the Sgr B region show
moderate, uniform subsidences, and we subtracted them by the linear fitting of
the data. The half-power-beam-width (HPBW), spatial grid, and velocity
resolution are 30″, 7$\farcs$2, and 1 km s-1, respectively.

As a shock tracer, we use SiO($J$=5–4) data that was also obtained by Ginsburg
et al. (2016). We subtracted moderate, uniform subsidences from these spectra,
by linear fitting of the data. The HPBW, spatial grid, and velocity resolution
for the data are also 30″, 7$\farcs$2, and 1 km s-1, respectively.

In order to calculate the intensity ratio between the $J$ = 2–1 and 1–0
transitions of C18O, we used the C18O($J$=1–0) data obtained by Tokuyama et
al. (2019) with the Nobeyama 45 m telescope. These spectra also show moderate,
uniform subsidences, and we subtracted them by linear fitting of the data. The
HPBW, spatial grid, and velocity resolution for the data are 21″, 7$\farcs$5,
and 2 km s-1, respectively.

###  3.2 Other Wavelengths

We mainly utilize the 90 cm radio-continuum data obtained by LaRosa et al.
(2000) as an indicator of star formation activity. In order to trace cold
dust, we utilize the FIR (250, 350, and 500 $\mu$m) data obtained by the
Spectral and Photometric Imaging Receiver installed on Herschel111Herschel is
an ESA space observatory with science instruments provided by European-led
Principal Investigator consortia and with important participation from NASA.
(Pilbratt et al., 2010; Griffin et al., 2010). We also use mid-infrared (MIR;
24, 8, and 3.6 $\mu$m) data obtained with the Multiband Imaging Photometer
(MIPS) and Infrared Array Camera installed on Spitzer (Rieke et al., 2004;
Fazio et al., 2004)

##  4 Results

###  4.1 Spatial Distribution of C18O($J$=2–1) toward the Sgr B region

Figure 1: Velocity channel distributions of C18O($J$=2–1) toward the Sgr B
region, including Sgr B1 and Sgr B2, obtained with APEX, overlaid with thick
black contours outlining the 90 cm radio-continuum emission. The red circles
indicate the positions of Sgr B2(N), (M), and (S), respectively. Known CO
features (the Shell, the Hole, and the Clump) are labeled.

Figure 1 shows the velocity channel distribution of C18O($J$=2–1) toward the
Sgr B region. Molecular clouds with velocities out of the figure are located
in the foreground or in the so-called “expanding molecular ring” (Kaifu et
al., 1972), and are not associated with Sgr B1 and Sgr B2 (see also Figure 4
of Vastel et al., 2002). The positions of the active star-forming regions (Sgr
B1, Sgr B2) and star-forming cores (Sgr B2(N), Sgr B2(M), and Sgr B2(S)) are
indicated in the top-left panel. The black contours indicate the boundaries of
the active star-forming regions obtained from the 90 cm radio-continuum data
(LaRosa et al., 2000). As seen in Figure 1, C18O($J$=2–1) traces the diffuse
gas well, and it peaks at $V_{LSR}$ $\sim 60$ km s-1 toward Sgr B2(M), meaning
that there is no significant self-absorption, even toward Sgr B2(M). At
$V_{LSR}$ = 20–40 km s-1, 40–50 km s-1, and 60–80 km s-1, the Shell appears as
a $\sim$30 pc cavity, the Hole is slightly north of Sgr B2, and the Clump is
located near the star-forming cores, as proposed by Hasegawa et al. (1994).
The figure shows that the Clump is obviously associated with the star-forming
cores. This has already been reported in the literature (e.g., Tsuboi et al.,
1999). The shapes of the Shell and the black contours agree well with each
other. The filamentary clouds elongated in the direction of Galactic longitude
in $V_{LSR}$ $\sim$80–90 km s-1 are probably located in another arm (arm II;
Sofue, 1995) or stream (Kruijssen et al., 2015).

###  4.2 Application of the CCC Identification Methodology

In this section, we use the identification methodology developed by Fukui et
al. (2018a) and Enokiya et al. (2021a) to investigate the possibility of a CCC
in the Sgr B region. As described in Section 2.2, we use the C18O($J$=2–1)
line as the best tracer for the methodology in this region.

Figure 2: Panels (a) and (b): intensity-weighted first-moment and second-
moment maps of C18O($J$=2–1) in 10 $\leq$ $V_{LSR}$ $\leq$ 80 km s-1. The
white circles indicate the positions of Sgr B2(N), (M), and (S). The thin and
thick black contours indicate the boundaries of the first-moment = 60.0 km s-1
and the radio-continuum emission at 90 cm, respectively. Panel (c):
longitude–velocity diagram for C18O($J$=2–1) toward the Sgr B region. The
black lines indicate the velocity range used for the moment calculations.

Following Enokiya et al. (2021a), we first define the velocity components in
this region by analyzing the moments of C18O($J$=2–1). Figures 2(a) and 2(b)
show the distributions of the intensity-weighted mean velocity and the
standard deviation of the velocity, respectively. We used only voxels with at
least a 5 $\sigma$ significance level for this analysis, in order to reduce
noise fluctuations. Figure 2(c) shows a longitude-velocity diagram toward the
Sgr B region. The associated clouds reported in previous works are seen as
major bright features in 10 $\leq$ $V_{LSR}$$\leq$ 80 km s-1, indicated by the
black lines. We use this velocity range to calculate the moments. Figure 2(a)
shows that redshifted features (from the interior to the thin black contour:
$V_{LSR}$ $\geq$ 60 km s-1) are seen toward the star-forming cores in Sgr B2.
This component contains the Clump (Hasegawa et al., 1994). We hereafter call
this velocity component the “60 km s-1 component”. The exterior of the black
contour shows velocity gradients ranging from yellow to blue in Figure 2a. We
hereafter call this velocity feature the “40 km s-1 component”. The eastern
contact line of the 60 km s-1 component has relatively larger line widths,
colored light green (Figure 2(b)).

In accordance with Enokiya et al. (2021a), we determine $V_{\mathrm{center}}$
and $dV$, respectively, as the averages of velocity of the first and second
moments in the given area, and we take the representative velocity range to be
$V_{\mathrm{center}}$ $\pm$ $dV$. We use the interior and exterior of the thin
black contour (first moment = 60.0 km s-1) in Figure 2(a) as the 60 km s-1 and
40 km s-1 component areas, respectively. We have confirmed that the difference
between these areas does not significantly change the results shown below. We
list the representative velocity range, $V_{\mathrm{center}}$, and $dV$ for
each velocity component in Table 1.

Table 1: Physical Parameters of the Velocity Components

Name | Representative velocity range | $V_{\mathrm{center}}$ | $dV$ |  $N_{\mathrm{H_{2}}}$ (typical/peak) | Mass
---|---|---|---|---|---
|  (km s-1) |  (km s-1) |  (km s-1) |  ($\times$1023 cm-2) |  ($\times$106 M⊙)
40 km s-1 component | 29 – 53 | 41.6 | 12.2 |  $\sim$0.5 / 1.7 | 1.4
60 km s-1 component | 53 – 74 | 63.8 | 10.4 |  $\sim$1 / 6.6 | 1.1

Note. — Column 1: names of the components, column 2: velocity ranges; column
3: peak velocities derived from the moment 1 map; column 4: velocity line
widths derived from the moment 2 map; column 5: typical and peak molecular
column densities toward each component; column 6: molecular masses.

Figure 3: Panel (a): integrated intensity distributions of the 40 km s-1
component (blue color) and 60 km s-1 component (orange contours) in
C18O($J$=2–1). Panels (b)and (c): velocity–latitude diagram and
longitude–velocity diagram of C18O($J$=2–1) toward the Sgr B region. The
velocity integration ranges for the 40 km s-1 and 60 km s-1 components are
indicated by the transparent blue and orange belts, respectively.

Figures 3(a)–3(c) show the integrated intensity distributions, the
velocity–latitude diagram, and the longitude–velocity diagram for the 40 km
s-1 and 60 km s-1 components in C18O($J$=2–1). The two velocity components
exhibit the typical signature of a CCC: a complementary distribution in space
(Figure 3(a)). The star-forming cores, indicated as the red circles, are
located at the intensity peak of the higher–column density cloud, the 60 km
s-1 component. Furthermore, there are rotated, V-shaped features in the
position–velocity diagrams. All of these features are typical signatures of a
CCC. We also investigated the SiO emissions and the line-intensity ratios
between C18O($J$=2–1) and C18O($J$=1–0) (R${}_{2-1/1-0}^{18}$) toward this
region (see the Appendix, Figures 9 and 10). From Figure 4 of Nishimura et al.
(2015), the typical value of R${}_{2-1/1-0}^{18}$ and those in the star-
forming regions in Orion are $\sim$0.5 and 1.0–1.2, respectively. In the Sgr B
region, a typical value of R${}_{2-1/1-0}^{18}$ is $\sim$0.8; it is higher at
the boundary of the Shell and the boundary of the Hole (1.0–1.6) and highest
at the Clump (1.6–2.8). The SiO($J$=5–4) emissions are detected only in the
areas where R${}_{2-1/1-0}^{18}$ is higher.

###  4.3 Physical Parameters of the Two Components

Here, we derive physical parameters, such as column densities and masses, for
the 40 km s-1 and 60 km s-1 components, assuming local thermodynamic
equilibrium (LTE). We assume that the excitation temperature is
$T_{\mathrm{ex}}$ = 30 K, which is similar to the temperatures of previous
studies (Hasegawa et al., 1994; Morris et al., 1983). Although the
temperatures may vary with position or component, the differences in
$T_{\mathrm{ex}}$ do not significantly affect the results.

First, we define the equivalent brightness temperature $J(T)$ as

$J(T)=\frac{h\nu}{\mbox{$k_{\mathrm{B}}$}}~{}[~{}\exp(\frac{h\nu}{\mbox{$k_{\mathrm{B}}$}T-1})~{}]^{-1}.$
(1)

where $h$, kB, and $\nu$ are Planck’s constant, Boltzmann’s constant, and the
frequency, respectively. A variation of the radiative transfer equation gives
the optical depth $\tau_{\nu}$ as

$\tau(v)=-\ln(~{}1~{}-~{}\frac{\mbox{$T_{\mathrm{obs}}$}}{J(\mbox{$T_{\mathrm{ex}}$})~{}-~{}J(\mbox{$T_{\mathrm{bg}}$})}~{}).$
(2)

where $T_{\mathrm{obs}}$ and $T_{\mathrm{bg}}$ are the observed intensity on
the $T_{\mathrm{mb}}$ scale and the temperature of the background (= 3 K),
respectively. The column density $N$ of the molecules can be calculated as
follows:

$N=\sum_{V}~{}\tau(v)~{}\Delta
v~{}\frac{3\mbox{$k_{\mathrm{B}}$}\mbox{$T_{\mathrm{ex}}$}}{4\pi^{3}\nu\mu^{2}}~{}\exp(\frac{h\nu
J}{2\mbox{$k_{\mathrm{B}}$}\mbox{$T_{\mathrm{ex}}$}})~{}\frac{1}{1-\exp({-\frac{h\nu}{\mbox{$k_{\mathrm{B}}$}\mbox{$T_{\mathrm{ex}}$}}})},$
(3)

where $V$, $\Delta v$, $\mu$, and $J$ are the velocity of the cloud, the
velocity resolution of the data, the electric-dipole moment of the molecule,
and the rotational transition level, respectively. Substituting
$k_{\mathrm{B}}$ = 1.38 $\times~{}10^{-16}$ (erg/K), $T_{\mathrm{ex}}$ = 30
(K), $\nu$ = 2.196 $\times~{}10^{11}$ (Hz), $\mu$ = 1.10 $\times~{}10^{-19}$
(esu cm), $h$ = 6.63 $\times~{}10^{-27}$ (erg s), and $J$ = 1 into Equation
(3) gives

$\mbox{$N_{\mathrm{C^{18}O}}$}=1.518\times
10^{16}~{}\sum_{V}~{}\tau(v)~{}\Delta v.$ (4)

The relation between $N_{\mathrm{C^{18}O}}$ and $N_{\mathrm{H_{2}}}$ was
determined by Frerking et al. (1982) to be

$\mbox{$N_{\mathrm{H_{2}}}$}=(\frac{\mbox{$N_{\mathrm{C^{18}O}}$}}{1.7\times
10^{14}}+1.3)\times 10^{21}.$ (5)

Using Equations (4) and (5), we find the typical / peak column densities of
the 40 km s-1 and 60 km s-1 components to be 0.5 / 1.7 $\times$ 1024 and 1 /
6.6 $\times$ 1023 cm-2, respectively.

The molecular mass is determined from following equation:

$M=\mbox{$\mu_{\mathrm{m}}$}\mbox{$m_{\mathrm{p}}$}\sum_{i}~{}[d^{2}\Omega\mbox{$N_{\mathrm{H_{2}}}$}_{,i}],$
(6)

where $\mu_{\mathrm{m}}$, $m_{\mathrm{p}}$, $d$, $\Omega$ and
$\mbox{$N_{\mathrm{H_{2}}}$}_{,i}$ are the mean molecular weight, proton mass,
distance, solid angle subtended by a pixel, and column density of molecular
hydrogen for the $i$th pixel, respectively. We assume a helium abundance of
20$\%$, which corresponds to $\mu_{\mathrm{m}}$ = 2.8. If we take $d$ = 8300
pc, the masses of the 40 km s-1 and 60 km s-1 components are $\sim$1.4 and
$\sim$1.1 $\times$ 106 M⊙, respectively. These physical parameters are listed
in Table 1.

###  4.4 Multiwavelength View of the Sgr B Region

Above, we have defined two velocity components (the 40 km s-1 component and 60
km s-1 component), and we have shown that a CCC model is applicable in this
region. In this section, we present a multiwavelength view of the region and
compare it with the results obtained in the previous section.

Figure 4: Panels (a) and (b): integrated intensity distributions of the 40 km
s-1 component and the 60 km s-1 component in C18O($J$=2–1), obtained with
APEX. The red circles indicate the positions of Sgr B2(N), (M), and (S), from
north to south. The thick black contours indicate the outline of the radio-
continuum emission at 90 cm. Panels (c) and (d): the cyan contours obtained
from the distributions of C18O($J$=2–1) in Figures 4(a), and 4(b),
respectively are superposed on a three-color composite image obtained with
Spitzer. The red, green, and blue correspond to 24, 8, and 3.6 $\mu$m,
respectively. Panels (e) and (f): the same contours of Figures 4(b), and 4(c),
respectively, are superposed on a three-color composite image obtained with
Herschel. Here, the red, green, and blue correspond 500, 350, and 250 $\mu$m,
respectively.

Figure 4(a) shows the spatial distribution of the 40 km s-1 component,
overlaid with the black contours of the boundary of the star-forming region
traced by the 90 cm radio-continuum emission. The representative velocity
ranges defined in Section 4.2 are used as the integration ranges. Figures 4(c)
and 4(e) are superpositions of contours of the contours of the 40 km s-1
component on the MIR and FIR images obtained with Spitzer and Herschel,
respectively. The FIR composite image shows only a single burgundy color. This
likely indicates that the properties of the cold dust emitting the FIR are the
same throughout this region, except for the white area, where the radiation is
very strong and saturated. The cold dust is seen as dark clouds in the MIR,
which means that a bright MIR source radiates from behind. The major radiative
process in the MIR is irradiation with UV photons from high-mass stars, so the
MIR is another indicator of star-forming activity, in addition to the radio-
continuum. It is clear that the bright MIR regions coincide well, excluding
the vicinity of the star-forming cores, with the radio-continuum boundary (the
black contour in Figure 4(a)). Thus, the bright background sources are Sgr B1
and Sgr B2.

The distribution of the 40 km s-1 component in the dense-gas tracer,
C18O($J$=2–1), is very similar to that of the cold dust or the MIR dark
clouds. These three therefore originate from the same object at different
wavelengths, and they are located in front of Sgr B1 and Sgr B2. At the
northwest edge of Sgr B1, the integrated intensity of C18O($J$=2–1) is
enhanced and well delineate the boundary of the radio-continuum emission
(Figure 4(a)); thus, the 40 km s-1 component is likely interacting with Sgr
B1. These are the dense clouds labeled “cloud d” and “cloud e” by Longmore et
al. (2013).

Figure 4(b) also shows the spatial distribution of the 60 km s-1 component in
C18O($J$=2–1), overlaid with the black contours of the boundary of the star-
forming region traced by the 90 cm radio-continuum emission. Figures 4(d) and
4(f) are similar superpositions of the contours of the 60 km s-1 component on
the MIR and FIR images obtained with Spitzer and Herschel, respectively. The
60 km s-1 component coincides with the strong FIR peak and the star-forming
cores in Sgr B2. This result clearly shows that the star-forming cores of Sgr
B2 are still at a very young stage of cluster formation, and are deeply
embedded in the high-density cloud with $N_{\mathrm{H_{2}}}$ $\geq$ 1023 cm-2,
as already reported in ealier studies. The 60 km s-1 component is also dark in
the MIR image, although this region is very bright in the radio-continuum.
Thus, the 60 km s-1 component is located a bit closer to the Sun than Sgr B2.
The darkness in the MIR image also indicates that Sgr B2 is a much younger
system than Sgr B1.

###  4.5 Summary of Observational Insights

Here, we summarize the observational aspects of the Sgr B region and the
results of our analyses.

  1. i)

Following Enokiya et al. (2021a), we identified two velocity features (the 40
km s-1 and 60 km s-1 components) associated with Sgr B1 and Sgr B2 by using
moment maps of C18O($J$=2–1).

  2. ii)

These two components clearly show all the signatures of a CCC proposed by
Fukui et al. (2018a): the complementary spatial distribution between the two
components, the V-shaped structure formed by the two colliding clouds in the
position–velocity diagrams, and the coincidence between the location of the
intensity peak of the denser component and the high-mass star-forming region.
We also detected SiO emission, which accompanies CCCs in the GC (Enokiya et
al., 2021a).

  3. iii)

Using the LTE approximation, we estimated the physical parameters of the two
velocity components with C18O($J$=2–1). The peak column densities and masses
of the 40 km s-1/60 km s-1 components are $\sim$2 $\times$ 1023 cm-2 / $\sim$7
$\times$ 1023 cm-2, and 1.4 $\times$ 106 M⊙ / 1.1 $\times$ 106 M⊙,
respectively.

  4. iv)

By comparing the CO distributions and other wavelengths, we found that the 40
km s-1 component, which is observed as an IR dark cloud, is located in front
of Sgr B1 and Sgr B2. The 60 km s-1 component is a high-density dust clump,
with $N_{\mathrm{H_{2}}}$ $\geq$ 1023 cm-2, which includes star-forming cores.
It is at a very young stage of cluster formation, and is located in front of
Sgr B2.

##  5 Discussion

In the previous Section, we applied the CCC identification methodology to the
Sgr B region and found clear signatures of a CCC. In this section, we
investigate further observational evidence for a CCC and discuss the CCC
scenario, including star-forming activities.

###  5.1 Comparison with Numerical Simulations

Figure 5: Panel (a): longitude-latitude-velocity surface plot of the Sgr B
region in C18O($J$=2–1). Panel (b): ppv surface plot of a numerical simulation
of a CCC performed by Takahira et al. (2014) and Haworth et al. (2015).

Figures 5(a) and 5(b), respectively, show a lngitude-latitude-velocity surface
plot of C18O($J$=2–1) in the Sgr B region and a position-position-velocity
(ppv) surface plot in the 12CO($J$=1–0) synthetic observational data obtained
by Haworth et al. (2015) (for a head-on collision with a relative velocity of
10 km s-1, observed from the angle of the relative motion). The color in
Figure 5(a) corresponds to velocity. Both ppv plots show similar, cone-shaped
structures. Henshaw et al. (2016) found a similar structure by using HNCO data
obtained with the Mopra telescope (Jones et al., 2012). Note that these
surface density plots were produced by velocities derived from the first
moment (intensity-weighted mean velocity), and hence the velocity of the
smaller cloud in the plot reflects the velocities of both the smaller and
larger clouds, and does not directly present its original velocity. Thus, the
quantitative comparison between Figure 5(a) and 5(b) does not make sense.

The conical structures seen in Figure 5 typically result from a collisional
interaction between two clouds, due to the exchange of the clouds’ momenta,
which is observed as the V shape in the position–velocity diagram (See Section
2.1). This structure is easily weakened by projection effects (see Fukui et
al., 2018a, for details). The observation shows a clear conical structure, but
the apex is shifted slightly from the center, in the direction of positive
Galactic longitude, suggesting that the collision may not be head-on and that
the 60 km s-1 component has a relative velocity toward positive Galactic
longitudes (see Fukui et al., 2018a, for details).

The part of the larger cloud that surrounds the cone shape, but that does not
participate in any collisional interactions, preserves the larger cloud’s
original velocity structures, as illustrated by the light and dark blue in
Figure 5(b). In the observational data, the velocity gradient of the cone
extends from the eastern boundary of Sgr B2 to the western boundary of Sgr B1.
Thus, if we adopt the CCC model, a comprehensive CCC probably occurred between
the two clouds, as shown in light blue and red in Figure 5(a). These two
clouds correspond to the Shell and the Clump. Sgr B1 is located in the
interior of the Shell, and thus it is possible that not only Sgr B2, but also
Sgr B1, were triggered by the CCC. Larger velocity widths are observed at the
collisional interface (that is, at the base of the cone structure in the
synthetic observations), and the enhanced line widths and SiO emissions
detected along the inner boundary of the Shell lend support to a comprehensive
collision at the 30 pc scale.

Figure 4(c) shows that the 40 km s-1 component, which appears dark, intercepts
the MIR emission from the background source, Sgr B1. In spite of this, the
boundary of the MIR emission from Sgr B1 coincides with the optically thin
radio-continuum emission, which is not intercepted by the foreground cloud of
the 40 km s-1 component (see Figure 4(a)). This means that the star-forming
activities of Sgr B1 only spread toward the region where the CO intensity of
the 40 km s-1 component is depressed. This is clear evidence that a past
collision at the position of Sgr B1 hollowed out the 40 km s-1 component and
created the hole, and that the collision triggered high-mass star formation in
Sgr B1. Conversely, if we abandon the CCC model, the coincidence between the
CO and radio-continuum toward Sgr B1 is rather unnatural.

###  5.2 CCC Scenario

Figure 6: Upper panel: schematic image of the time evolution of the CCC in the
Sgr B region (top view). Lower panel: schematic image of the time evolution of
the longitude–velocity diagram.

From the above discussion, we here propose the following plausible collision
scenario: two redshifted clouds—or one large redshifted cloud with two density
peaks—approached a blueshifted cloud (the leftmost panel of Figure 6). About 3
Myr ago, the two velocity features started to collide at the position of Sgr
B1, and formed Sgr B1 first (the middle panel of Figure 6). The region
involved in this collision then extended from west to east (from Sgr B1 to Sgr
B2). A recent collision at the positions of Sgr B2(N), (M), and (S) triggered
star-forming activities there (the rightmost panel of Figure 6). The
blueshifted cloud is likely to be the 40 km s-1 component. If we assume the
age of Sgr B1 to be 3 Myr, then, by dividing the distance between Sgr B2 and
Sgr B1 by the age, we obtain the speed at which the collision spread along the
Galactic longitude as 30 pc / 3 Myr = 10 km s-1.

We suggest two possible candidates for the redshifted cloud(s). One
possibility is that the clouds are located in arm II at $V_{LSR}$ $\sim$90 km
s-1, as shown in Figures 1 and 2(c). These figures show that the 60 km s-1
component connects to arm II in space and velocity. In this case, the velocity
difference of $\sim$30 km s-1 between the 60 km s-1 component and arm II can
be explained by the exchange of momenta between the 40 km s-1 component and
arm II (90 km s-1). The other possibility is a cloud with an initial velocity
of $\sim$60 km s-1. If this cloud had a very high-density, compared to that of
the 40 km s-1 component, the cloud might maintain its initial velocity even
after the collision. However, there is so far no clear evidence for the
existence of clouds with 60 km s-1 in this direction, and furthermore, the 60
km s-1 component has apparently extended gas to the higher velocity (see e.g.,
Figure 3(b)), then the first case might be more plausible.

According to Fukui et al. (2018a) two velocity components are not always
observable in a CCC, due to projection effects and so on. Fukui et al. (2018b)
present a case of a CCC with a single-velocity component in the star cluster
GM24. Therefore, even though the conical structure in the ppv plot consists of
a single-peak velocity profile, this model is still possible.

Figure 7: Panel (a): close-up view of integrated intensity distribution of the
40 km s-1 component (blue color) and the 60 km s-1 component (orange
contours). The red circles indicate the positions of Sgr B2(N), (M), and (S).
Panel (b): the same as Figure 7(a), but the contours are displaced 4 pixels
(1.16 pc) to the west and 2 pixel (0.58 pc) to the north. The length and the
direction of the displacement is indicated by the yellow arrow in Figure 7(a).

Figure 7(a) shows distributions of the 40 km s-1 component (blue contours) and
the 60 km s-1 component (orange contours). The distributions of the hole in
the 40 km s-1 component and the 60 km s-1 component display a slight
displacement. Following Enokiya et al. (2021a), and by using an algorithm
developed by Fujita et al. (2021), we estimate the displacement to be
$\sim$1.3 pc to the northwest, which maximizes the complementarity between the
two components. The yellow arrow in Figure 7(a) indicates the derived
displacement vector. Figure 7(b), which shows the contours displaced by that
arrow overlaid on the 40 km s-1 component, exhibits perfect coincidence
between the two components. Three active star-forming cores are located at the
center of the 60 km s-1 component. If we assume that the LOS relative velocity
and the relative velocity in the direction of Galactic longitude are same,
then by dividing the displacement by the relative velocity, we estimate the
duration time of the collision at the star-forming cores to be 1 pc/30 km s-1
$\simeq$ 4 $\times$ 104 yr. This value is comparable to the typical age of
hypercompact Hii regions (103–104 yr, e.g., Qin et al., 2008). From numerical
MHD simulations, Fukui et al. (2021) reported that molecular flows colliding
at a relative velocity of 20 km s-1 and with densities of 1000 cm-3 generate
massive clumps at the collision interface, with mass accretion rates larger
than 10-4 M⊙ yr-1. Since the 60 km s-1 component is denser than their
simulations, a CCC is a viable process for creating the star-forming cores.

###  5.3 Origin of the Large Velocity Difference between the Two Clouds

The driving force responsible for producing the large relative velocity of a
few tens of km s-1 between two massive clouds is still unclear, but we propose
two possibilities. One is a bar-potential-driven model. As Binney et al.
(1991) originally proposed, two families of x1 and x2 orbits can be produced
in the central 1 kpc of a galaxy under certain bar-potential conditions. The
probability of a CCC increases at the intersections of these orbits (or arms).
According to the x1–x2 orbits dynamical model, as in Binney et al. (1991),
Bally et al. (2010), and Molinari et al. (2011), the Sgr B region is located
at one of the intersections.

Based on hydrodynamic simulations, Sormani et al. (2019) discussed the origins
of extended velocity features (EVFs), and suggested that EVFs are consequently
obtained as results from two types of cloud collisions, overshot type or
connection type, accelerated by the bar potential. The overshot type requires
an interloper, which is overshoot from the dust lane. The redshifted clouds
likely belong to arm II (see Section 5.2), and have not been overshoot; thus,
the other cloud—i.e., the blueshifted cloud, or, in other words, the 40 km s-1
component—should be overshot from a dust lane resulting from bar-potential
acceleration, if this model is true. Although the mass and size of the
overshoot cloud is quite large (see Table1), and thus it is apparently
unnatural, it still remains a possible scenario. If the Sgr B region is a
connection type EVF, it needs to be a connecting point for at least two
streams. As discussed at the beginning of this section, this scenario is also
possible.

Another possibility is a magnetically driven model. The GC has strong magnetic
fields at least 50 $\mu$G (Crocker et al., 2010), and such conditions can
easily excite Alfvén velocities of the order of a few tens of km s-1 (see
Suzuki et al., 2015). Some authors have also reported observational evidence
of loop-like gas streams perpendicular to the Galactic plane produced by
magnetic buoyancy (Fukui et al., 2006; Enokiya et al., 2014). The probability
of a CCC increases at the footpoints of these loops. Enokiya et al. (2021a)
discovered an enhancement in the probability of a CCC at the footpoints of the
loops. They suggested that the enhancement is caused by rear end collisions
originated from the stagnation of the flow at the footpoints. According to the
open-orbit dynamical model, the orbit has warps in $z$-direction. Although
this warp has so far not been clearly explained, if we assume that the warp is
reflecting the shape of toroidal magnetic field line, the Sgr B region is
located at the foot point of the loop (i. e., the base of the warp). Thus, the
collision at the Sgr B region can be interpreted as a rear end collision
resulting from the gravitationally accelerated cloud sliding down along the
loop.

###  5.4 The Star Formation Triggered by the CCC Event

Figure 8: Panel (a): H2 column-density map of the molecular gas associated
with the Sgr B region estimated from the C18O($J$=2–1) emission in 20 km s-1
$\leq$ $V_{LSR}$ $\leq$ 90 km s-1. The contours show the distribution of the
radio-continuum emission at 90 cm. The cyan circles indicate the positions of
Sgr B2(N), (M), and (S). The integration latitude ranges for Figures 8(b) and
8(c) are indicated by the black lines and dashed black lines, respectively.
Panel (b): the integrated column density (blue) and integrated 90 cm flux
(black) plots toward the star-forming regions in the Sgr B region. Panel (c):
the same plots for the limited latitude range focused on the Sgr B2 star-
forming cores.

The image and contours in Figure 8(a) show the distributions of the column
densities ($N_{\mathrm{H_{2}}}$) derived by using the LTE approximation and
the radio-continuum emission, respectively. We adopted $T_{\mathrm{ex}}$ = 30
K and the velocity range 20 $\leq$ $V_{LSR}$ $\leq$ 90 km s-1 for the
calculation. Figure 8(a) clearly shows that the column density is typically
$\sim$1023 cm-2 and increases up to $\sim$1024 cm-2 toward the star-forming
cores, and it decreases to less than $\sim$1022 cm-2 toward Sgr B1. We assume
here that the radio-continuum emission is optically thin, in which case it is
a good tracer of star formation activity. Figures 8(a) shows that the
correlations between gas and star-formation activity get better/worse toward
Sgr B2/Sgr B1.

The plots of the integrated $N_{\mathrm{H_{2}}}$ and radio-continuum flux
versus the Galactic longitude shown in Figures 8(b) and 8(c) delineate the
above picture more quantitatively. The integrated latitude ranges of Figure
8(b) and 8(c) correspond to the two star-forming regions and the star-forming
cores, respectively. Figure 8(b) shows that the star-forming activities in Sgr
B1 and Sgr B2 are fairly comparable, although the mass of the molecular gas
concentrated in Sgr B2 is one order of magnitude higher than in Sgr B1. This
may mean that 90$\%$ of the molecular gas in Sgr B1 has been consumed to form
stars or has been dissipated by ionization or the CCC. As discussed in Section
5.1, the shape of the cavity in the 40 km s-1 component and Sgr B1 are very
similar, although they have slightly different LOS distances. Therefore, the
collision may have hollowed out the 40 km s-1 component, creating the cavity
and triggering high-mass stars, thus dissipating the natal cloud. Figure 8c
shows that $\sim$50$\%$ of the star-forming activity and the molecular gas in
Sgr B2 are contained in a small volume. On the other hand, the molecular gas
in Sgr B1 is significantly decreased. Again, this depression is likely caused
by the collision and ionization.

###  5.5 Position of the Sgr B Region among Other CCC Regions

Enokiya et al. (2021a) collected all the literature reporting the signatures
of CCCs, and they carried out a statistical analysis of the data. For more
than 50 Galactic objects, they found correlations between the number of O- and
early B-type stars and the peak column densities of the colliding clouds, as
well as between the relative velocities and peak column densities of the
colliding clouds (see Figure 9 of Enokiya et al., 2021a). The CCC in the Sgr B
region shows the following relative velocity, peak column density, and number
of high-mass stars (candidates): 22 km s-1, $\sim$1024, and a few hundreds,
respectively, in the star forming cores. These physical parameters mostly
satisfy the correlations found by Enokiya et al. (2021a) for the most massive
cases.

The correlation plots indicate that only a CCC with a lower column density
triggers the formation of a normal star cluster, that a moderate column
density triggers a superstar cluster, and that a high column density may
trigger a starburst. The threshold column densities for triggering a normal
cluster, a superstar cluster, and a starburst are $\sim$1021.5, $\sim$1023,
and $\sim$1024 cm-2, respectively. Some authors have suggested that a
collision between two large-scale Hi flows can trigger a massive cluster
containing M∗ of a few hundred thousands M⊙, such as R136 in the Large
Magellanic Cloud (Fukui et al., 2017; Tachihara et al., 2018). Detailed
molecular observations of the star-forming cores in Sgr B2 with ALMA may
reveal the CCC star-forming process at a very young stage, and link the
triggering mechanism with the active star formation seen in the external
galaxies.

##  6 Summary

We have used a recently developed CCC identification methodology to
investigate the possibility of a CCC in the Sgr B region, by using datasets at
various wavelengths and fully sampled molecular-line data. Our major findings
are summarized as follows.

  1. 1.

We identified two CO components, with velocities of $\sim$42 (the 40 km s-1
component) and $\sim$64 km s-1 (the 60 km s-1 component) by using
C18O($J$=2–1) as the optically thin, dense-gas tracer. These velocity
components form a cone-like structure covering both Sgr B1 and Sgr B2 in the
ppv plot.

  2. 2.

From comparisons with other wavelengths, we identified the locations of the
velocity components and the star-forming regions in the Sgr B region as
follows: Sgr B1 and Sgr B2 are located farthest from the Sun, while the 40 km
s-1 component is nearest, and intercepts the light from Sgr B1 and Sgr B2. The
star-forming cores in Sgr B2—namely Sgr B2(N), (M), and (S), accompanied by
the 60 km s-1 component—are located nearer to the remaining region of Sgr B2.
According to previous work, Sgr B2 is likely to be located nearer to the Sun
than Sgr B1. Although the LOS distances to Sgr B1 and the 40 km s-1 component
are likely to be slightly different, the distributions of these two regions
show an interaction between them.

  3. 3.

Two clouds satisfy all the signatures of a CCC, as suggested by Fukui et al.
(2018a). If we apply the CCC model, Sgr B1 and Sgr B2 are explained as results
of contiguous collisions between the 40 km s-1 and redshifted clouds. One of
the cores of the redshifted cloud, with a column density of 1023-24 cm-2,
collided first and formed Sgr B1 a few Myr ago, while a very recent collision
of another core of the redshifted cloud a few $\times$ 104 yr ago triggered
active star formation in the star-forming cores in Sgr B2, which include a few
hundred high-mass stars. The candidates for the redshifted cloud are either a
cloud with a velocity of $\sim$60 km s-1 or the so-called arm II.

  4. 4.

The CCC in the Sgr B region satisfies the correlation among the relative
velocity, peak column density, and the number of O- and early B-type stars
found in Enokiya et al. (2021a) in the most active star formation case.

Although we do not have clear evidence for this, the mechanism responsible for
driving the large relative velocity of the clouds may either be an encounter
between arms or streams, or else a magnetic instability, such as the Parker
instability. Higher-resolution observations will be required in order to
determine the detailed mechanism responsible for triggering star formation in
Sgr B2.

We thank the anonymous referee(s) for helpful comments that improved the
manuscript. RE is grateful to Dr. T. Hasegawa for helpful discussions and
comments on this work. The authors would like to thank Enago (www.enago.jp)
for the English language review. This work is based [in part] on observations
made with the Spitzer Space Telescope, which is operated by the Jet Propulsion
Laboratory, California Institute of Technology under a contract with NASA.
This work was financially supported Grants-in-Aid for Scientific Research
(KAKENHI) of the Japanese society for the Promotion for Science (JSPS; grant
number society for 20K14520).

##  Appendix A Appendix information

Here we show the distributions of the shocked gas (Figure 9) and the relative
excitation states (Figure 10) toward the Sgr B1 and Sgr B2 regions. Figures 9
and 10 show the velocity channel distributions of SiO($J$=5–4) and
R${}_{2-1/1-0}^{18}$, respectively.

Figure 9: Velocity channel distributions of SiO($J$=5–4) toward the Sgr B
region, obtained with APEX, overlaid with the thick black contours of the 90
cm radio-continuum emission. The red circles indicate the positions of Sgr
B2(N), (M), and (S). Figure 10: Velocity channel distributions of the
intensity ratios of C18O($J$=2–1) and C18O($J$=1–0) toward the Sgr B region,
overlaid with the thick black contours of the 90 cm radio-continuum emission.
The red circles indicate the positions of Sgr B2(N), (M), and (S).

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