11institutetext: University of Chinese Academy of Sciences, Beijing 100049,
People’s Republic of China
11email<EMAIL_ADDRESS><EMAIL_ADDRESS>22institutetext: I.
Physikalisches Institut, Universität zu Köln, Zülpicher Str. 77, D-50937 Köln,
Germany 33institutetext: Department of Astronomy, University of Florida, PO
Box 112055, USA 44institutetext: NRAO, 520 Edgemont Rd, Charlottesville, VA,
USA 55institutetext: Agnes Scott College, 141 E. College Ave., Decatur, GA
30030, USA 66institutetext: Max Planck Institute for Extraterrestrial Physics,
Giessenbachstrasse 1, D-85748 Garching, Germany 77institutetext: Max Planck
Institute for Radio Astronomy, Auf dem Hügel 69, D-53121 Bonn, Germany

# The physical and chemical structure of Sagittarius B2

VI. UCHii regions in Sgr B2

F. Meng 1122 Á. Sánchez-Monge 22 P. Schilke 22 A. Ginsburg 33 C. DePree 4455
N. Budaiev 33 D. Jeff 33 A. Schmiedeke 66 A. Schwörer 22 V. S. Veena 77 Th.
Möller 22

(Received ; accepted )

###### Abstract

Context.

The giant molecular cloud Sagittarius B2 (hereafter Sgr B2) is the most
massive region with ongoing high-mass star formation in the Galaxy. Two ultra-
compact Hii (UCHii ) regions were identified in Sgr B2’s central hot cores,
Sgr B2(M) and Sgr B2(N).

Aims. Our aim is to characterize the properties of the Hii regions in the
entire Sgr B2 cloud. Comparing the Hii regions and the dust cores, we aim to
depict the evolutionary stages of different parts of Sgr B2.

Methods. We use the Very Large Array in its A, CnB, and D configurations, and
in the frequency band C ($\sim$6 GHz) to observe the whole Sgr B2 complex.
Using ancillary VLA data at 22.4 GHz and ALMA data at 96 GHz, we calculated
the physical parameters of the UCHii regions and their dense gas environment.

Results. We identify 54 UCHii regions in the 6 GHz image, 39 of which are also
detected at 22.4 GHz. Eight of the 54 UCHii regions are newly discovered. The
UCHii regions have radii between $0.006\,{\rm pc}$ and $0.04\,{\rm pc}$, and
have emission measure between $10^{6}\,{\rm pc\,cm^{-6}}$ and $10^{9}\,{\rm
pc\,cm^{-6}}$. The UCHii regions are ionized by stars of types from B0.5 to
O6. We found a typical gas density of $\sim 10^{6}-10^{9}\,{\rm cm^{-3}}$
around the UCHii regions. The pressure of the UCHii regions and the dense gas
surrounding them are comparable. The expansion timescale of these UCHii
regions is determined to be $\sim 10^{4}-10^{5}\,{\rm yr}$. The percentage of
the dust cores that are associated with Hii regions are 33%, 73%, 4%, and 1%
for Sgr B2(N), Sgr B2(M), Sgr B2(S), and Sgr B2(DS), respectively. Two-thirds
of the dust cores in Sgr B2(DS) are associated with outflows.

Conclusions. The electron densities of the UCHii regions we identified are in
agreement with that of typical UCHii regions, while the radii are smaller than
those of the typical UCHii regions. The dust cores in Sgr B2(M) are more
evolved than in Sgr B2(N). The dust cores in Sgr B2(DS) are younger than in
Sgr B2(M) or Sgr B2(N).

###### Key Words.:

Stars: formation – Stars: massive – Radio continuum: ISM – Radio lines: ISM –
ISM: clouds – ISM: individual objects: Sgr B2

##  1 Introduction

The giant molecular cloud Sagittarius B2 (Sgr B2) is the most massive ($\sim
10^{7}\,M_{\odot}$) region with ongoing high-mass star formation in the Galaxy
(see e.g., Goldsmith et al., 1990). Sgr B2 has a higher density
($>10^{5}\rm\,cm^{-3}$) and dust temperature ($\gtrsim$50–70 K) compared to
other star forming regions in the Galactic plane (see e.g., Ginsburg et al.,
2016; Schmiedeke et al., 2016; Sánchez-Monge et al., 2017). Sgr B2 is located
at a distance of $8.34\pm 0.16$ pc, and only $\sim$100 pc in projection from
the Galactic center (Reid et al., 2014)111 A new distance to the Galactic
center of $8.127\pm 0.031$ kpc has been measured (Gravity Collaboration et
al., 2018). For consistency with the papers published within the same series
of studies of Sgr B2, we use the distance reported by Reid et al. (2014). .
These features make Sgr B2 an excellent case to study high-mass star formation
in an extreme high-pressure environment. Such an environment resembles nearby
starburst galaxies (Leroy et al., 2018). Understanding the structure of the
Sgr B2 molecular cloud complex is necessary to comprehend the most massive
star forming region in our Galaxy, which at the same time provides a unique
opportunity to study in detail the nearest counterpart of the extreme
environments that dominate star formation in the Universe (see, e.g.,
Kruijssen & Longmore, 2013; Henshaw et al., 2022). This paper continues our
series of studies on Sgr B2 (Schmiedeke et al., 2016; Sánchez-Monge et al.,
2017; Pols et al., 2018; Schwörer et al., 2019; Meng et al., 2019). In Meng et
al. (2019) we presented the observations of Sgr B2(DS), which is a part of Sgr
B2 giant cloud, and analyzed the physical properties of the non-thermal
emission within it. In this work we study the whole Sgr B2 region.

In the central $2\,{\rm pc}$ of Sgr B2 there are the two well-known and well-
studied hot cores Sgr B2(N) and Sgr B2(M) (see, e.g., Schmiedeke et al., 2016;
Sánchez-Monge et al., 2017), which contain at least 70 high-mass stars with
spectral types from O5 to B0 (see, e.g., Gaume et al., 1995; De Pree et al.,
1998, 2014). Surrounding the two hot cores, there is a larger envelope
(hereafter _the envelope_) with a radius of 20 pc that contains more than 99%
of the total mass of Sgr B2 (Schmiedeke et al., 2016). Along with the active
high-mass star forming activity discovered in Sgr B2(N) and Sgr B2(M), hints
of star formation happening in the envelope are also revealed. Ginsburg et al.
(2018), with ALMA at 3 mm, revealed 271 high-mass protostellar cores
distributed throughout the entire Sgr B2 region, including the envelope. The
luminosities of these dust cores suggest that they must contain objects with
stellar masses higher than 8 $M_{\odot}$.

Due to the high extinction in the infrared bands toward Sgr B2 (see Meng et
al., 2019), there is no direct evidence of the existence of high-mass stars
embedded in the dust cores detected by Ginsburg et al. (2018). However, since
high-mass stars ionize the neutral material surrounding them, the presence and
properties of the associated Hii regions reflect the evolutionary stages of
these dust cores (see, e.g., González-Avilés et al., 2005; Breen et al.,
2010). Additionally, since the free-free emission from Hii regions may extend
from centimeter to millimeter wavelengths in the spectral domain (see, e.g.,
Sánchez-Monge et al., 2013), measuring the luminosities of the associated Hii
regions can help us better constrain the luminosities of the dust cores.
Therefore, to further characterize the evolutionary stages and physical
properties of these dust cores, we investigate the possible Hii regions
associated with them.

The Hii regions in Sgr B2 were targeted by several previous studies. Mehringer
et al. (1993) observed the entire Sgr B2 with VLA in the 20, 6, and 3.6 cm
bands and identified 15 Hii regions. The resolutions range from $\sim$20″ to
$\sim$3″ when the wavelength changes from 20 to 3.6 cm, which correspond to
the range 0.8–0.12 pc. The 15 Hii regions, except two unresolved cases, all
have sizes $>2$″. Since the 271 dust cores may contain newly formed high-mass
stars (Ginsburg et al., 2018), the associated Hii regions may be ultra-compact
Hii (UCHii) and hyper-compact Hii (HCHii) regions, which typically have sizes
from $\sim 0.03$ pc to $\sim 0.1$ pc (see, e.g., Kurtz, 2002; González-Avilés
et al., 2005; Kurtz, 2005; Breen et al., 2010). Thus, the resolution of
Mehringer et al. (1993) is not sufficient to resolve the UCHii and HCHii
regions. Gaume & Claussen (1990); Gaume et al. (1995) observed Sgr B2(N) and
Sgr B2(M) at 7 mm and 1.3 cm and achieved resolutions of 0.065″ and 0.25″,
respectively. Rolffs et al. (2011) observed Sgr B2(N) and Sgr B2(M) in 40 GHz
with a resolution of 0.1″. Unfortunately, these high-resolution observations
do not cover the entire envelope. For example, Sgr B2(DS), where $\sim$80 of
dust cores reside (see, e.g., Ginsburg et al., 2018; Meng et al., 2019), is
left out of these observations. LaRosa et al. (2000); Law et al. (2008a, b)
also observed the entire Sgr B2 at centimeter wavelengths, but the resolutions
were not high enough to study the UCHii and HCHii regions.

In this paper we present Very Large Array (VLA) observations of the entire Sgr
B2 cloud in the frequency regime 4–8 GHz, with configurations A, BnC, and D.
The high resolution ($\lesssim 0.01$ pc) and large spatial coverage ($\sim 20$
pc) of our data sets make a systemic and complete study of the UCHii and HCHii
regions in Sgr B2 possible. We also include analysis of the 3 mm image
(Ginsburg et al., 2018) as well as the newly acquired SiO (5–4) data, both of
which were observed with the Atacama Large Millimeter/submillimeter Array
(ALMA). Thus, we can disentangle the contributions of ionized gas and dust at
millimeter wavelengths and better constrain the evolutionary stages of the
dust cores.

This paper is organized as follows. In Sect. 2 we describe the observations
and the data reduction process. In Sect. 3 we present the results. In Sect. 4
we discuss the results. Finally, we summarize this paper in Sect. 6.

##  2 Observations and data reduction

We used the VLA in its A, CnB, and D configurations to observe the entire Sgr
B2 complex in frequency band C (4–8 GHz). In the following we call this band 6
GHz. The observations with the CnB and D configurations were described in Meng
et al. (2019). The observations with the A configuration were conducted from
October 1 to 12, 2016 (project 16B-031, PI: F. Meng). We used 64 spectral
windows with a bandwidth of 128 MHz each. Mosaic mode was used, with ten
pointings for C band. The primary beam of each pointing is 7.5′. Quasar 3C286
was used as the flux and bandpass calibrator, the SED of which is
$S_{\nu}=5.059\pm 0.021\ {\rm Jy}\times(S/8.435\ {\rm GHz})^{-0.46}$ from 0.5
to 50 GHz (Perley & Butler, 2013). Quasar J1820-2528, whose flux is 1.3 Jy in
the C Band, was used as phase calibrator. The calibration was done using the
standard VLA pipelines provided by the NRAO222The National Radio Astronomy
Observatory is a facility of the National Science Foundation operated under
cooperative agreement by Associated Universities, Inc..

Calibration and imaging were done in Common Astronomy Software Applications
(CASA 4.7.2 McMullin et al., 2007). The details of the data processing from
the CnB and D configurations are described in Meng et al. (2019). The A
configuration data were originally taken every 2 s. To shorten the time of
processing, we applied timebin in CASA to the measurement sets, which averaged
the data taken within 10 s into one data point. All the pointings of the
mosaic in each band were primary beam corrected and the mosaic was imaged
using the CASA task tclean. With a robust factor of 0, the image of C band has
a synthesized beam of 0.62″$\times$0.28″, with a position angle ${\rm(PA)}$ of
$7.27^{\circ}$. The PA is defined positive north to east. To mitigate the
spatial filtering effect of the A configuration image, we applied the feather
algorithm to combine the images from A configuration and from CnB and D
configurations. The combined image of the three configurations is shown in
Fig. 1. The combined image has a resolution identical to that of the A
configuration images, while being sensitive to spatial scales up to $\sim
240$″. The sensitivities of the observations are described in Sect. 3.

Ancillary data include the 22.4 GHz data, the 3 mm continuum, and SiO (5–4)
data. The observations of the 22.4 GHz data are described in Gaume et al.
(1995), with a resolution of 0.27″$\times$0.23″ (${\rm PA}=70^{\circ}$). The
root mean square (RMS) noise is 0.38 mJy/beam. The spatial coverage of the
22.4 GHz data is shown as the blue dashed box in Fig. 1. The 96 GHz continuum
data covers the frequency range from 89.5 to 103.3 GHz. The image in 96 GHz
has a resolution of 0.54″$\times$0.46″ (${\rm PA}=68.31^{\circ}$), the
observational details of which are described in Ginsburg et al. (2018). Unlike
the 22.4 GHz image, the 96 GHz image covers the entire area shown in Fig. 1.
The 22.4 GHz observations were performed 27 years prior to the 6 GHz and 96
GHz observations. If we assume that Sgr B2 has a proper motion of $\sim 50$
${\rm km\ s^{-1}}$ (Henshaw et al., 2016), the corresponding position shift in
27 years is 0.04″, which is only $\sim$10% of the beam sizes of our images. In
addition, we did not find any visible shift in the positions of the compact
sources positions between the 6 and 22.4 GHz images. Thus, we do not take the
astrometry difference into account when analyzing the 6 and 22.4 GHz data in
the following sections. The SiO (5–4) emission was observed with ALMA (Project
2017.1.00114.S, P.I. A. Ginsburg) and has a resolution of 0.35″$\times$ 0.24″,
with PA of $-80^{\circ}$, and spectral resolution of 1.35 km s-1. For the
details of the observation and data reduction, see Jeff et al. (in prep.). The
typical RMS of the SiO image is 0.9 mJy/beam (0.3 K). The observation covers
Sgr B2(S) and the eastern part of Sgr B2(DS).

Figure 1: Sources identified in the 6 cm image. Those objects shown as white
and blue circles have radii of $r_{\rm obs6}$ (see Table 2). Notable Hii
regions are indicated. The Hii regions that are masked out for core
identification are shown as white contours. The beam size is given in the
lower left corner of the main panel $(\theta_{\rm maj},\theta_{\rm min},{\rm
PA})$. The spatial coverage of the 22.4 GHz image is shown as a blue dashed
box.

##  3 Results

In this section we present the image of Sgr B2 at 6 GHz and the UCHii regions
identified in it. For the Hii regions, we calculated their actual sizes and
physical properties using the observations at 6 GHz and at 22.4 GHz. With the
96 GHz data we characterized the properties of the dense gas that the UCHii
regions reside in.

###  3.1 Observed parameters of the 6 GHz sources

Figure 1 displays the image of Sgr B2 at 6 GHz, where the known large-scale
Hii regions are denoted N, M, S, AA, DS, and V, following the nomenclature of
Mehringer et al. (1992, 1993); Ginsburg et al. (2018); Meng et al. (2019). The
maps at 22.4 GHz and 96 GHz are presented by Gaume et al. (1995) and Ginsburg
et al. (2018), respectively. Our aim is to study UCHii regions and since the
22.4 GHz image has non-complete spatial coverage, we performed compact source
identification at the 6 GHz image only. We used the automatic source
extracting algorithm SExtractor (Bertin & Arnouts, 1996), which allows us to
identify the location of bright compact sources throughout the map. We note
that this algorithm also includes a large fraction of artificial sources due
to some large-scale artifacts visible in the final image (see the zoomed-in
images in Fig. 1). The variable noise (see Fig. 13) also has an effect on the
exclusion of certain sources with the automatic algorithms. After cross-
checking the automatically produced lists of sources at 6 GHz and 22.4 GHz, we
made the final catalog of continuum sources at 6 GHz by excluding or adding
sources by visual inspection (in a similar way to the approach followed in
Ginsburg et al. 2018 for the ALMA 96 GHz data). Since we focus on compact
sources, extended emission larger than 5″ or 0.2 pc (see contours in Fig. 1)
are excluded in the catalog and the following analysis. In total, 54 compact
sources are identified throughout the entire Sgr B2 cloud (see Table 2), 8 of
which are newly identified. Of these 54 cores, 8 are identified in Sgr B2(N),
40 in Sgr B2(M), 2 in Sgr B2(S), and 1 in Sgr B2(DS). All 54 compact sources
are covered by the 22.4 GHz image, except core 1.

For each of the 54 compact sources we define a minimal circle that can include
as much as the total flux density of it at 6 GHz.333 This is done by expanding
the circle as long as the flux $S_{\rm total}$ within the circle increases
monotonically. When $S_{\rm total}$ remains constant or (due to the “negative
bowl” in the artifacts) decreases, the corresponding circle is called the
minimum circle.  The observed radius of the core and the flux density within
the core are denoted $r_{\rm obs6}$ and $S_{\rm 6}$, respectively. For these
54 sources we followed the same photometry procedure at the 22.4 GHz and 96
GHz bands, and obtained $r_{\rm obs22}$, $S_{\rm 22}$, $r_{\rm obs96}$, and
$S_{\rm 96}$ that correspond to the observed radii and flux densities. The
observed radius and flux density in three bands of all the 54 compact sources
are listed in Table 2.

To match the sources across the three images (6, 22.4, and 96 GHz), we defined
that if two sources in two images have a distance (distance between their
centers) shorter than either of their radii, they are associated with each
other. Afterward, we manually adjust few of the matched sources in the crowed
region in Sgr B2(M). For all the compact sources at 6 GHz, emission in the 96
GHz band is also detected. Among the 15 sources that have no $S_{\rm 22}$, 14
are without reliable 22.4 GHz detection ($S_{22}<3{\rm RMS}$), and one (#1 in
Tab. 2) is not spatially covered by the 22.4 GHz image. All the 22.4 GHz
sources444 In Gaume et al. (1995), extended Hii regions are also included in
the catalog. We only compare the compact sources ($r<5$″) in their catalog
with ours. reported by Gaume et al. (1995) are associated with the 6 GHz
compact sources.

We compared the spatial association of the 54 sources with the 271 dust cores
identified at 96 GHz by Ginsburg et al. (2018). Although there is no one-to-
one correspondence between the 271 dust cores and the associated dust emission
of the 54 compact sources, we can still conclude that at least 217 (80%) dust
cores are not associated with compact radio sources that are detectable with
our sensitivity.

Figure 2: Probability distribution of $S_{\rm 6}$ (pink), plotted as a kernel
density estimation. Mean and median values are indicated. The distributions of
$S_{\rm 22}$ (green) and $S_{\rm 96}$ (blue) are also shown.

Most of the compact sources have $r_{\rm obs6}$ from 0.5″ to 1″, which is
comparable to the beam size at 6 GHz. Even though the 22.4 GHz image has
higher resolution ($\sim 0.25^{\prime\prime}$), the measured $r_{\rm obs22}$
are still not significantly larger than the beam size at 22.4 GHz. Therefore,
most of the compact sources are not well resolved; in other words, $r_{\rm
obs6}$ or $r_{\rm obs22}$ cannot accurately represent the actual radii of most
of the compact sources. The 96 GHz compact sources, as indicated by Ginsburg
et al. (2018), are also not resolved. The flux densities of the compact
sources are distributed in a wide range; $S_{\rm 6}$ ranges from $\sim 3$ mJy
to $\sim 300$ mJy at 6 GHz, as shown in Fig. 2. The mean and median of $S_{\rm
6}$ are 37.3 mJy and 16.2 mJy, respectively.

To test whether the nature of the emission at 6 and 22.4 GHz is free-free
emission, we interpolated $S_{\rm 6}$ and $S_{\rm 22}$ to derive a spectral
index $\alpha_{\rm 6-22}$ for each of the 39 sources that have 22.4 GHz
detection, see the last column of Table 2. All 39 sources have $\alpha_{\rm
6-22}>-0.1$, which suggest that their emission at centimeter wavelengths is
very probably dominated by thermal free-free emission from ionized gas (see,
e.g., Sánchez-Monge et al., 2013). For the 15 sources without 22.4 GHz
detections, we find that their $S_{\rm 6}$ are not higher than the values of
the other sources. If we treat their emission at 6 GHz as free-free as well
555For single-dish images, the RMS around a certain null detection spot could
be used to derive an upper limit of the possible signal, which in this paper
can be translated into an upper limit of $\alpha_{\rm 6-22}$; however, because
the artifacts in the 22.4 GHz image contain negative bowls of interferometric
images, we do not trust the upper limits derived from the 22.4 GHz RMS map. ,
the derived physical parameters are also within the same ranges as the other
sources (see Section 3.2). Thus, we treat all 54 sources as thermal free-free
sources in the following analysis.

###  3.2 Physical parameters of the Hii regions

As we show in Sect. 3.1, most of the compact sources at 6 and 22.4 GHz are not
well resolved. Therefore, we first need to determine their actual sizes. The
observed flux density ($S_{\rm 6}$ and $S_{\rm 22}$) of an Hii region is
related to its actual size ($r_{\rm calc}$, which is called calculated radius
in this paper), its electron temperature ($T_{\rm e}$), and its emission
measure (EM) as (see, e.g., Wilson et al., 2013)

$\frac{S_{\nu}}{\rm mJy}=8.183\times 10^{-4}\left(\frac{r_{\rm calc}}{\rm
arcsec}\right)^{2}\left(\frac{\nu}{\rm GHz}\right)^{2}\left(\frac{T_{\rm
e}}{\rm K}\right)(1-e^{-\tau_{\nu}}),$ (1)

where the optical depth $\tau_{\nu}$ is

$\tau_{\nu}=8.235\times 10^{-2}\left(\frac{\nu}{\rm
GHz}\right)^{-2.1}\left(\frac{T_{\rm e}}{\rm K}\right)^{-1.35}\left(\frac{\rm
EM}{\rm pc\,cm^{-6}}\right).$ (2)

In our case the frequency $\nu$ is 6 GHz or 22.4 GHz. We assume that $T_{\rm
e}$ is $10^{4}$ K, following the values given by Mehringer et al. (1993) for
Sgr B2(N) and Sgr B2(M), considering that most of the compact sources are in
these two regions. Thus, with the two independent flux density measurements,
$S_{6}$ and $S_{22}$, we solve Eq. 1 and Eq. 2 to obtain $r_{\rm calc}$ and EM
simultaneously for the 39 sources that are detected at both 6 and 22.4 GHz.
For the 15 compact sources that are only detected at 6 GHz, we deconvolved the
beam size from $r_{\rm obs6}$ to estimate $r_{\rm calc}$, as $r_{\rm
calc}^{2}+r_{\rm beam}^{2}=r_{\rm obs6}^{2}$, and $r_{\rm
beam}^{2}=0.62^{\prime\prime}\times 0.28^{\prime\prime}$ for 6 GHz. Then we
calculated EM using the estimated $r_{\rm calc}$. The $r_{\rm calc}$ and EM of
all the compact sources are listed in Table 3; those derived from the
deconvolved $r_{\rm calc}$ are flagged with and asterisk (*) in Table 3.

Figure 3: Relation between $r_{\rm obs6}$ and $r_{\rm calc}$ of the 39 compact
sources with 22.4 GHz detection. The dashed curve indicates the deconvolution
relationship $r_{\rm calc}^{2}+r_{\rm beam}^{2}=r_{\rm obs6}^{2}$, where
$r_{\rm beam}$ is the effective radius of the beam.  Figure 4: Distribution of
EM. The mean and median values are indicated.  Figure 5: Distribution of
$n_{\rm e}$. The mean and median values, which are both $\sim 10^{5}\,{\rm
cm^{-3}}$, are indicated.

For all the compact sources that are associated with the 22.4 GHz detection,
we plot the $r_{\rm obs6}-r_{\rm calc}$ diagram (see Fig. 3). We can see that
in general $r_{\rm calc}$ follows the trend of deconvolution (i.e., $r_{\rm
calc}^{2}=r_{\rm obs6}^{2}-0.62^{\prime\prime}\times 0.28^{\prime\prime}$),
which suggests that the method we used to calculate $r_{\rm calc}$, although
without measuring the compact source size in the image, is in general
consistent with the size we observed. The deviation between the $r_{\rm calc}$
from the dashed curve in Fig. 3 is likely due to some uncertainties in the
observations and flux measurements. For example, the determination of $r_{\rm
obs6}$ is made by including as much of the flux as possible, which might be
larger than the actual observed size. In addition, due to the artifacts in the
image, we cannot ideally define the boundaries of the compact sources, but
only approximate the compact source as a circle with $r=r_{\rm obs}$.

Figure 6: Distribution of $\dot{N}_{\rm Ly}$. The corresponding spectral type
of the ionizing stars (assuming that the Hii regions are ionized by single
stars) are indicated.

The probability distribution of EM of all the 54 compact sources is shown in
Fig. 4. The EM of these compact sources ranges from $\sim 2\times 10^{6}\,{\rm
pc\,cm^{-6}}$ to $\sim 3\times 10^{9}\,{\rm pc\,cm^{-6}}$. Only eight compact
sources have ${\rm EM}<10^{7}\,{\rm pc\,cm^{-6}}$, which means that most of
the compact sources are consistent with the parameters derived for UCHii
regions (see Kurtz, 2002). From EM we estimated the electron density $n_{\rm
e}$ as $n_{\rm e}=[{\rm EM}/(2r_{\rm calc})]^{1/2}$. The distribution of
$n_{\rm e}$ is shown in Fig. 5. The $n_{\rm e}$ of the 54 Hii regions are from
$\sim 4\times 10^{3}\,{\rm cm^{-3}}$ to $\sim 5\times 10^{5}\,{\rm cm^{-3}}$
with a mean value of $\sim 10^{5}\,{\rm cm^{-3}}$.

Knowing EM, we calculated the flux of Lyman continuum photons ($\dot{N}_{\rm
Ly}$) that are needed to ionize these UCHii regions following Eq. C.18 in
Schmiedeke et al. (2016). The derived $\dot{N}_{\rm Ly}$ ranges from
$10^{46}\,{\rm s^{-1}}$ to $10^{49}\,{\rm s^{-1}}$ (see Fig. 6). If we assume
that these cores are ionized by single stars, these UCHii regions are ionized
by stars from spectral types B0.5 to O6 (Panagia, 1973).

In this section we describe a method that is independent of the observed
source size to derive the radius ($r_{\rm calc}$) and emission measure (EM) of
39 sources detected at 6 and 22.4 GHz. To verify this a method, we compared
the $r_{\rm calc}$, EM, and $\dot{N}_{\rm Ly}$ values of these 39 sources with
the values presented by Gaume et al. (1995), namely $r_{\rm G95}$, $\rm
EM_{G95}$, and $\dot{N}_{\rm Ly\,G95}$. Of the 39 sources, 33 are coincident
with the UCHii regions reported by Gaume et al. (1995). Twenty of the 33
sources have $r_{\rm calc}>r_{\rm G95}$ (see left panel of Fig. 7). The order
of magnitude of EM of the UCHii regions is consistent in these two studies. Of
the 33 sources, 23 have EM larger than $\rm EM_{G95}$ (see middle panel of
Fig. 7). Most of the sources have $\dot{N}_{\rm Ly}$ similar to $\dot{N}_{\rm
Ly\,G95}$ (see right panel of Fig. 7), except source 44. The difference
between $\dot{N}_{\rm Ly}$ and $\dot{N}_{\rm Ly\,G95}$ of source 44 is
possibly caused by the time-domain flickering of the radio emission in SgrB2
(see the last paragraph of this section).

Figure 7: Comparison between $r_{\rm calc}$ (_left_), EM (_middle_), and
$\dot{N}_{\rm Ly}$ (_right_) of 33 sources and the corresponding parameters
found by Gaume et al. (1995).  Figure 8: Diagram of $r_{\rm calc}$ and $n_{\rm
e}$. Previously derived $r_{\rm calc}$–$n_{\rm e}$ relationships for UCHii
regions from Garay & Lizano (1999), $n_{\rm e}=780\times(2r)^{-1.19}$, and
from Kim & Koo (2001), $n_{\rm e}=790\times(2r)^{-0.99}$, are plotted as
dashed lines. The green shading indicates the typical $n_{\rm e}$ of UCHii
regions, which is from $\sim 10^{4}\,{\rm cm^{-3}}$ to $\sim 10^{6}\,{\rm
cm^{-3}}$, and the typical $r$ of UCHii regions, which is $\sim 0.015$ to
$\sim 0.05$ pc (Kurtz, 2005).

In Fig. 8 we plot the diagram of $r_{\rm calc}$ and $n_{\rm e}$. It is worth
noting that all of the 39 Hii regions have typical $n_{\rm e}$ of UCHii
regions (i.e., $\sim 10^{4}\,{\rm cm^{-3}}$ to $\sim 10^{6}\,{\rm cm^{-3}}$;
Kurtz 2005). However, 16 of these have $r$ smaller than that of typical UCHii
regions ($\sim 0.015$ to $\sim 0.05$ pc). Although some of the sources have
sizes similar to HCHii regions, considering the low density of our Hii regions
compared to typical HCHii regions, and following the previous nomenclature of
most of the sources (e.g., Gaume et al., 1995), we still call all the Hii
regions in this study UCHii regions.

Recently, Rivera-Soto et al. (2020) identified similar Hii regions in W 51
that have $n_{\rm e}\sim 10^{4}-10^{5}\,{\rm cm^{-3}}$ and $2\times r\sim
10^{-3}-10^{-2}\,{\rm pc}$. They categorized these sources as HCHii regions,
but they are similar to smaller UCHii regions and are ionized by early B-type
stars, which is in agreement with the spectral type of the ionizing stars for
most of our UCHii regions (see Fig. 6). Unlike the HCHii regions in W51, our
UCHii regions do not exactly follow the $r-n_{\rm e}$ relationships of UCHii
regions proposed by Garay & Lizano (1999) and Kim & Koo (2001), and have
higher $n_{\rm e}$ (by a factor of $\sim 2$) than the predicted values (see
Fig. 8). Such a discrepancy might be due to the neutral gas surrounding our
UCHii regions that is denser ($n_{\rm H_{2}}\gtrsim 10^{6}\,{\rm cm^{-3}}$)
than typical molecular cores ($n_{\rm H_{2}}\sim 10^{4}-10^{5}\,{\rm
cm^{-3}}$; see, e.g., Bergin & Tafalla 2007). In Sect. 3.3 a detailed analysis
of the gas density surrounding the UCHii regions is presented.

Since we cannot resolve the detailed morphologies of the UCHii regions, we
neglect the possible inhomogeneity of the Hii regions. The various
morphologies of UCHii regions result in modified SEDs other than that
described by Eq. 1 and Eq. 2, (see, e.g., Keto, 2003; Keto et al., 2008).
Additionally, if there are accretion flows to the Hii regions, flickering of
the flux on a timescale of $\sim 100$ yr may occur, which is observed by De
Pree et al. (2014) and modeled by Peters et al. (2010b, c). Since the 6 GHz
observation (2013) was performed 23 years after the 22.4 GHz observations
(1989), the effects of flickering could be present in some of these cores.
Hence, the simultaneous use of 6 GHz and 22.4 GHz fluxes may not be
appropriate for the analysis of some sources. However, De Pree et al. (2014)
reported the flickering of 4 out of 41 sources in Sgr B2 within a similar time
range (1989–2014). Such a rarity of cases (10%) suggests that the flickering
may not significantly alter our statistical results.

###  3.3 Dense gas environment

We extrapolated the Hii regions’ SED (Eq. 1) to subtract the free-free
contribution from $S_{96}$ to get the emission purely from dust, which is
denoted $S_{\rm dust}$. Twelve UCHii regions have $S_{\rm dust}$ below
$3\times$rms, while the other 42 UCHii regions have detectable dust emission
after subtraction of the free-free emission. Using $S_{\rm dust}$, we
evaluated the dust properties in the vicinity of the 42 UCHii regions.

Following Ossenkopf & Henning (1994), we calculated the dust column density
$N_{\rm d}$ as

$S_{\rm dust}=\frac{2h\nu^{3}}{c^{2}}\frac{1}{e^{\frac{h\nu}{k_{\rm B}T_{\rm
d}}}-1}\left(1-e^{-\kappa_{0}\left(\frac{\nu}{\nu_{0}}\right)^{\beta}N_{\rm
d}}\right)\frac{\pi r^{2}}{D^{2}},$ (3)

in which $T_{\rm d}$ is the dust temperature, $D$ is the distance of Sgr B2,
and $r$ is the radius of the dust core. Since most of the dust cores are not
resolved (Ginsburg et al., 2018), here we use $r=(r_{\rm obs96}^{2}-r_{\rm
beam}^{2})^{1/2}$, where $r_{\rm beam}=(0.54^{\prime\prime}\times
0.46^{\prime\prime})^{1/2}$ is the effective radius of the 96 GHz beam. The
dust parameters $\kappa_{0}$ and $\beta$ depend on the dust grain properties.
Due to the general high temperature and high density of neutral gas revealed
by previous studies (see, e.g., Huttemeister et al., 1993; Schmiedeke et al.,
2016), we assume that $T_{\rm d}=100$, $\kappa_{0}=2.631$, and $\beta=1.05$
for $\nu_{0}=100\,{\rm GHz}$. The assumption of $\kappa_{0}$ and $\beta$
corresponds to dust grains without ice mantles and with a volume density of
$10^{8}\,{\rm cm^{-3}}$. We assume a gas-to-dust mass ratio of 100 (Ott et
al., 2014; Giannetti et al., 2017) to estimate the gaseous mass $N_{\rm
H_{2}}$, $N_{\rm H_{2}}=100N_{\rm d}$.

From the column density of molecular gas, $N_{\rm H_{2}}$, we obtain the
volume density of molecular gas (H2),

$n_{\rm H_{2}}=\frac{3N_{\rm H_{2}}}{4r}.$ (4)

We list the $n_{\rm H_{2}}$ values at the position of the 42 UCHii regions
associated with dust emission in Table 3. The distribution of $n_{\rm H_{2}}$
is shown in Fig. 9. The $n_{\rm H_{2}}$ ranges from $\sim 10^{6}$ to $\sim
10^{8}\,{\rm cm^{-3}}$. Due to the artifacts and possible small size of the
dust core, even a dust core with enough volume density might not be detected
by the sensitivity of our data set. We used the method in Ginsburg et al.
(2018) to estimate the detection limit of the dust emission. Compared to the
typical $n_{\rm H_{2}}$ of the molecular clouds surrounding UCHii regions
($10^{5}\,{\rm cm^{-3}}$) (e.g., Wood & Churchwell, 1989), our UCHii regions
reside in denser neutral gas. The median value obtained in this work is
similar to the value ($2\times 10^{7}\,{\rm cm^{-3}}$) measured in the central
region of Sgr B2(M) (de Pree et al., 1995; Huettemeister et al., 1995;
Sánchez-Monge et al., 2017), although our sources are also distributed in the
outskirts of Sgr B2(M).

Figure 9: Probability distribution of $n_{\rm H_{2}}$ of the dust surrounding
the UCHii regions that are associated with dust emission. Median and mean
values are indicated.

##  4 Analysis and discussion

###  4.1 Expansion and equilibrium

Figure 10:  Evolution sequence of the Hii regions. The circles are calculated
with the dust density derived from the 96 GHz image. The triangles are
calculated with the dust density of $2\times 10^{7}\,{\rm cm^{-3}}$. The color
of the markers indicates the expansion timescale.

Based on the physical properties of the Hii regions (see Sect. 3.2) and the
properties of their surrounding environment (see Sect. 3.3), we now
investigate the expansion time for these extremely compact and dense Hii
regions. The density of the molecular cloud that UCHii regions expand into
significantly affects the expansion rate (e.g., Wood & Churchwell, 1989; de
Pree et al., 1995; De Pree et al., 1998). For the simple case of a spherical
UCHii region ionized by the Lyman continuum flux of $\dot{N}_{\rm Ly}$ and
expanding in molecular gas with volume density ${n}_{\rm H_{2}}$ and electron
temperature of $10^{4}\,{\rm K}$, following de Pree et al. (1995); De Pree et
al. (1998) we calculate the initial Strömgren radius $r_{\rm i}$ as

$\frac{r_{\rm i}}{\rm pc}=1.99\times 10^{-2}\left(\frac{\dot{N}_{\rm Ly}}{\rm
10^{49}\,s^{-1}}\right)^{1/3}\left(\frac{{n}_{\rm H_{2}}}{\rm
10^{5}\,cm^{-3}}\right)^{-2/3}.$ (5)

For the 12 UCHii regions without detected associated dust emission ($S_{\rm
dust}$, which is $S_{96}$ subtracting the free-free component), we assumed a
uniform dust density of $2\times 10^{7}\,{\rm cm^{-3}}$ following De Pree et
al. (2015). This value is in agreement with the median of $n_{\rm H_{2}}$ of
the 42 sources that have physical $S_{\rm dust}$. Since we assumed a
relatively high gas density compared to typical dust cores (see, e.g., Bergin
& Tafalla, 2007), the nondetection of $S_{\rm dust}$ for 12 sources under such
an assumption is due to small $r$ but not low $n_{\rm H_{2}}$. Then we
calculated the expansion timescales of all 54 cores applying the expansion
equation by Spitzer (1968); Dyson & Williams (1980),

${}r_{\rm calc}=r_{\rm i}\left(1+\frac{7c_{\rm i}t}{4r_{\rm i}}\right)^{4/7},$
(6)

where $c_{\rm i}$ is the sound speed ($\sim 10\,{\rm km\,s^{-1}}$). The
expansion times, $t$, of the UCHii regions are listed in Table 3. In Fig. 10
the cores are color-coded according to their $t$. In ideal cases most of the
cores have an expansion time between $10^{4}$ yr and $10^{5}$ yr. The
expansion times in this section were dervied under the assumption of ideal
conditions (Spitzer, 1968). Such derived expansion times should be treated as
lower limits and the actual expansion time might be longer for to at least
three reasons. First. the expansion of UCHii regions can be halted due to the
pressure equilibrium between the ionized gas and the surrounding dense
molecular gas (e.g., De Pree et al., 1998). We assume that $T_{\rm e}$ of the
ionized gas is $10^{4}\,{\rm K}$, and the molecular temperature $T_{\rm
H_{2}}$ is from $50\,{\rm K}$ to $100\,{\rm K}$. The pressure equilibrium
condition ($2n_{\rm e}T_{\rm e}=n_{\rm H_{2}}T_{\rm H_{2}}$) can be expressed
as $2n_{\rm e}=5\times 10^{-3}n_{\rm H_{2}}$ and $2n_{\rm e}=10^{-2}n_{\rm
H_{2}}$ for $T_{\rm H_{2}}=50\,{\rm K}$ and $100\,{\rm K}$, respectively. In
dust cores with $n_{\rm H_{2}}\gtrsim\times 10^{7}\,{\rm cm^{-3}}$, the UCHii
regions are mostly in equilibrium with the neutral gas, evident from Fig. 11,
whereas for the lower $n_{\rm H_{2}}$ regime, the UCHii regions are in
expansion phase owing to ionized gas pressure exceeding the neutral gas
pressure. Second, in reality, dust absorption may reduce $r_{\rm i}$, and
therefore slows down the expansion (see, e.g., Wood & Churchwell, 1989; De
Pree et al., 1998). Third, accretion flow on to the central star will disrupt
the expansion process and cause a sudden decrease in the flux and size of the
UCHii region. (see, e.g., Fig. 10 in Peters et al., 2010a). With the current
data, we cannot quantify the effect of either of these mechanisms. Future
observations, for example high-resolution radio recombination line
observations, may provide constraints on the possible effect of accretion on
the expansion of the Hii regions.

Figure 11:  Diagram of $n_{\rm e}$ and $n_{\rm H_{2}}$. The Hii regions in Sgr
B2(M) are shown as circles, while those in Sgr B2(N) are shown as triangles.
The color indicates the size $r_{\rm calc}$ (see color scale at right).
Equilibrium conditions are plotted as solid and dashed lines for $T_{\rm
H_{2}}=100\,{\rm K}$ and $50\,{\rm K}$ and thus for $T_{\rm e}=100T_{\rm
H_{2}}$ and $200T_{\rm H_{2}}$.

###  4.2 Evolutionary stages

Table 1: Number of objects associated with dust cores

Region | N | M | S | DS | Resta | Total
---|---|---|---|---|---|---
Dust core | 24 | 55 | 46 | 46 | 100 | 271
Hii region | 8 | 40 | 2 | 1 | 3 | 54
Outflow | –b | –b | –b | 31 | 18b | 49b
CH3OH Maser | 2 | 2 | 0 | 0 | 5 | 9

a: Rest of the envelope.
b: Not fully covered by the SiO map.

The type of objects that are associated with dust cores suggests the
evolutionary stages of the dust cores in star formation activity. In Table 1
we summarize the type and number of objects associated with the dust cores in
the subregions Sgr B2(N), Sgr B2(M), Sgr B2(S), and Sgr B2(DS), and in the
rest of the envelope of Sgr B2.

Of all the 54 UCHii regions, 8 are in Sgr B2(N), 40 are in Sgr B2(M), 2 are in
Sgr B2(S), 1 is in Sgr B2(DS) and 3 are in the rest of the envelope. Although
the dust cores are distributed all over Sgr B2 (Ginsburg et al., 2018) and
more than 80 of them are associated with the large Hii region in Sgr B2(DS)
(Meng et al., 2019), the dust cores are rarely associated with any UCHii
regions outside of Sgr B2(N) and Sgr B2(M). The percentage of the dust cores
that are associated with UCHii regions are 33%, 73%, 4%, and 1% for Sgr B2(N),
Sgr B2(M), Sgr B2(S), and Sgr B2(DS), respectively, while for the remaining
part of the envelope the percentage is 3%.

For a dust core, association with the Hii region is a sign of that the core is
more evolved (see, e.g., Breen et al., 2010). Therefore, the cores in Sgr
B2(M) are the most evolved, while the cores in Sgr B2(S) and Sgr B2(DS) are
the least evolved. The evolutionary stages of the cores in Sgr B2(N) are
between those of Sgr B2(M) and Sgr B2(DS) or Sgr B2(S).

All the dust cores in Sgr B2(DS) appear to be pure dust666 One Hii region is
found associated with a core reported by Ginsburg et al. (2018), but the core
emission is pure free-free, which suggests that the core is a dust-free Hii
region. Therefore, no dust cores in Sgr B2(DS) are associated with Hii
regions.. With the sensitivity of our observation ($\sim 1$ mJy, for
unresolved sources), we constrain the age of the possibly undetectable
embedded Hii regions in the dust cores. We consider the $\dot{N}_{\rm Ly}$
from the central star are of $10^{46}$, $10^{47}$, $10^{48}$ ${\rm s^{-1}}$,
and dust densities are $10^{6}$ and $10^{8}$ ${\rm cm^{-3}}$, which are based
on the typical values of the detected UCHii regions. The evolution of $S_{6}$
of the Hii regions, based on the expansion model in Sect. 4.1 and Eq. 8 in (De
Pree et al., 1998), are displayed in Fig. 12. The temporal evolution of $r$
and EM leads to increase in $S_{\rm 6}$ (see Eq. 1). For the dust cores that
have $n_{\rm H_{2}}=10^{6}$ ${\rm cm^{-3}}$, the embedded Hii regions with
$\dot{N}_{\rm Ly}\geq 10^{47}$ ${\rm s^{-1}}$ will be observable $\sim 100$ yr
after their birth, which contradicts the nondetection of Hii regions in Sgr
B2(DS). For the dust cores that have $n_{\rm H_{2}}=10^{8}$ ${\rm cm^{-3}}$,
the embedded Hii regions with $\dot{N}_{\rm Ly}\geq 10^{47}$ ${\rm s^{-1}}$
will be observable $\sim 10^{3}$ yr after their birth, which suggests that it
is possible that the Hii regions are already formed but still too dim to be
detected. The Hii regions with $\dot{N}_{\rm Ly}\leq 10^{46}$ ${\rm s^{-1}}$
will be always undetectable with the current sensitivity. The sizes of the
embedded Hii regions will reach the typical size of the UCHii regions in this
study ($\sim 0.01\,{\rm pc}$) in $10^{4}$ years (de Pree et al., 1995, see
Fig. 1 in). Thus, it is possible that stars later than B1 have already formed
and ionized the interiors of the dust cores.

Figure 12: Modeled evolution of $S_{6}$ of Hii regions with time. The black
line denotes the sensitivity of our 6 GHz observation, which is $\sim 1$ mJy,
for unresolved sources.

Collimated outflows are the footprints of the very early stages of star
formation activity (see, e.g., Beuther et al., 2002). Outflows can be traced
by SiO emission (e.g., Schilke et al., 1997). In a recent project, the SiO
(5–4) emission was observed with ALMA (P.I. A. Ginsburg). For the details of
the observation and data reduction, see Jeff et. al. (in prep.). The
resolution is 0.35″ $\times$ 0.24″, with P.A. of $-80^{\circ}$. The spectral
resolution is 1.35 km s-1. The typical RMS of the image is 0.9 mJy/beam. The
observation covers Sgr B2(S) and the eastern part of Sgr B2(DS). The peak
intensity map of SiO (5–4) is shown in Figure 14. Due to the artifacts around
Sgr B2(S), we only analyzed the part of the image with declination
$<-$28:24:00. We generated the moment one map by masking out all the pixels
below $3\sigma$ (see Fig. 15). The average velocity difference between the
blueshifted and redshifted lobes is found to be $\sim 10$ ${\rm km\ s^{-1}}$ .
We visually matched the positions of the dust cores identified by Ginsburg et
al. (2018) and the SiO outflows. The dust cores that are 1) spatially
associated with an outflow, 2) covered by the SiO image not associated with an
outflow, and 3) not covered by the SiO image are indicated in Table 4.
Spatially, most of the outflows are associated with dust cores. On the other
hand, of all the 120 dust cores covered by the SiO image, 49 are identified as
associated with outflows. Particularly, among all the cores in Sgr B2(DS),
two-thirds are associated with outflows. Such a high fraction confirms that
the cores in Sgr B2(DS) are at their very early evolutionary stages.

CH3OH masers serve as additional probes to trace the star formation activity
(see, e.g., Breen et al., 2010). We cross-matched the CH3OH masers presented
by Caswell et al. (2010) with our dust cores777We searched for associated
masers in a circle with a radius of 1 arcsec, which is double the beam size of
the 96 GHz image.. In Table 1 the number of masers that are associated with
the dust cores that are in each region are listed. The cores in Sgr B2(M) and
in Sgr B2(N) are associated with two CH3OH masers. The cores in Sgr B2(DS)
have no associated masers. This lack of associated masers also suggests that
the cores in Sgr B2(DS) are less evolved than those in Sgr B2(N) and Sgr
B2(M).

It is worth noting that the dust cores in Sgr B2(DS) are distributed around
the large-scale Hii region in Sgr B2(DS) (See Fig. 2 in Meng et al., 2019).
The large-scale Hii region was ionized by a central O7 star and is still in
the expansion phase (Meng et al., 2019). The expansion time of this Hii region
is $\sim 10^{5}\,{\rm yr}$, which is estimated based on the size of the Hii
region (0.72 pc) and the sound speed of the ionized gas, $10\,{\rm
km\,s^{-1}}$. This timescale, compared to the expansion time of the Hii
regions in Sgr B2(M) and Sgr B2(N) suggest that the central O7 star possibly
formed before the stars in Sgr B2(M) and Sgr B2(N). Another possible scenario
is that the O7 star was ejected from a nearby star forming site (e.g., Sgr B1)
earlier than $10^{5}\,{\rm yr}$ ago, and Sgr B2 had no star forming activity
at that time. The newly formed dust cores in Sgr B2(DS), on the other hand,
are signs of star forming activity that possibly was triggered by the
expansion of the large-scale Hii region in Sgr B2(DS).

##  5 Summary

We observed the Sgr B2 complex with VLA A, BnC, and D array at 6 GHz and
identified 54 compact radios sources. We found that 39 of the 54 sources are
also detected in the 22.4 GHz band (Gaume et al., 1995). Our main results are
summarized as follows:

  * •

Using the 6 GHz data of all 54 sources, as well as the 22.4 GHz data of 39
sources, we calculated the EM, radius, electron densities, and the spectral
type of the ionizing stars of all the 54 UCHii regions. The UCHii regions have
radius between $6\times 10^{-3}\,{\rm pc}$ and $4\times 10^{-2}\,{\rm pc}$,
and have EM between $10^{6}\,{\rm pc\,cm^{-6}}$ and $10^{9}\,{\rm
pc\,cm^{-6}}$. We found that the electron densities of these UCHii regions are
in agreement with the values of typical UCHii regions, while the radii are
smaller than for the typical UCHii regions. We identified that the UCHii
regions are ionized by stars with spectral types ranging between B0.5 to O6.

  * •

Using the 96 GHz ALMA data, we characterized the dense gas environment where
the UCHii regions are located. We found a typical dense gas density of $\sim
10^{6}-10^{9}\,{\rm cm^{-3}}$ around the UCHii regions. Using Spitzer (1968),
we estimated the expansion timescale of the UCHii regions as $\sim
10^{4}-10^{5}\,{\rm yr}$. More than half of the UCHii regions are close to
equilibrium with the neutral gas; this means that the pressure of most UCHii
regions and the dense gas surrounding them are comparable. Due to the high
pressure of the neutral gas, some natal Hii regions might be optically thick
and cannot be detected with the sensitivity of this study. Instead, for the
lower $n_{\rm H_{2}}$ regime ($n_{\rm H_{2}}\lesssim\times 10^{7}\,{\rm
cm^{-3}}$) the ionized gas pressure exceeds that of the neutral gas and the
UCHii regions are still expanding.

  * •

The percentage of the dust cores that are associated with Hii regions are 33%,
73%, 4%, and 1% for Sgr B2(N), Sgr B2(M), Sgr B2(S), and Sgr B2(DS),
respectively. Among all the dust cores in Sgr B2(DS), two-thirds are
associated with outflows that are traced by SiO(5–4). The dust cores in both
of Sgr B2(M) and Sgr B2(N) are associated with two 6.7 GHz CH3OH masers, while
the dust cores in Sgr B2(DS) have no associated maser. Based on these
findings, we suggest that the dust cores in Sgr B2(M) are more evolved than
those in Sgr B2(N). The dust cores in Sgr B2(DS) are younger than those in Sgr
B2(M) or Sgr B2(N).

In this work, we calculated the physical properties of the UCHii regions and
their surrounding neutral gas in Sgr B2. We found that the pressure of the
UCHii regions and the dense gas surrounding them are comparable. We also
characterized the evolutionary stages of these UCHii regions and obtained
their minimum expansion timescales.

###### Acknowledgements.

FM, ASM, PS, ASchw research is carried out within the Collaborative Research
Centre 956, sub-projects A6 and C3, funded by the Deutsche
Forschungsgemeinschaft (DFG) - project ID 184018867. AG acknowledges support
from the NSF via AST 2008101 and CAREER 2142300. This paper makes use of the
following ALMA data: 2013.1.00269.S and 2017.1.00114.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. This research made use of
Astropy,888http://www.astropy.org a community-developed core Python package
for Astronomy (Astropy Collaboration et al., 2013, 2018).

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##  Appendix A RMS maps

The rms maps were generated using SExtractor (Bertin & Arnouts 1996). Although
the core identification was performed by eye, SExtractor produces rms maps as
side-products of automated core identification. The rms maps are shown in Fig.
13.

Figure 13: RMS maps of 6 GHz, 22.4 GHz, and 96 GHz. The resolutions are given
in Sect. 2. Each image has two profile cuts, in blue and red. To better show
the features and maxima of the rms of each image, the offset is different for
each of the three images. The offset of each profile increases from left to
right and from bottom to top. The profile is below each image.

##  Appendix B Tables

Table 2: Observed parameters

# | RA | DEC | $r_{\rm obs6}$ | $S_{6}$ | $r_{\rm obs22}$ | $S_{22}$ | $r_{\rm obs96}$ | $S_{96}$ | $\alpha_{6-22}$
---|---|---|---|---|---|---|---|---|---
| 17:47:- - |  $-$28:- - | arcsec | mJy | arcsec | mJy | arcsec | mJy |
1∗ | 23.355 | 25:33.95 | 0.70 | $2.4\pm 0.1$ | … | … | 0.35 | $0.3\pm 0.2$ | …
2a | 19.485 | 24:39.75 | 1.45 | $64.9\pm 0.3$ | … | … | 1.20 | $60.9\pm 0.2$ | …
3∗ | 20.235 | 23:44.95 | 0.60 | $7.0\pm 0.1$ | 0.60 | $14.7\pm 0.7$ | 0.60 | $7.9\pm 0.2$ | $0.56\pm 0.26$
4a | 20.428 | 23:44.25 | 2.10 | $363.1\pm 0.8$ | 1.15 | $820.5\pm 1.6$ | 1.45 | $1107.6\pm 0.5$ | $0.62\pm 0.27$
5a | 20.042 | 23:18.00 | 1.65 | $123.8\pm 0.2$ | 1.10 | $160.4\pm 0.9$ | 1.05 | $135.5\pm 0.2$ | $0.20\pm 0.09$
6a | 20.049 | 23:12.60 | 1.20 | $68.1\pm 0.2$ | 0.65 | $126.1\pm 0.7$ | 0.80 | $124.8\pm 0.2$ | $0.47\pm 0.21$
7b | 19.765 | 23:09.90 | 0.75 | $15.6\pm 0.1$ | 0.50 | $22.8\pm 0.7$ | 0.40 | $18.3\pm 0.2$ | $0.29\pm 0.13$
8∗ | 20.015 | 23:08.85 | 0.55 | $28.3\pm 0.3$ | 0.55 | $64.2\pm 0.7$ | 0.55 | $52.5\pm 0.6$ | $0.62\pm 0.27$
9a | 20.106 | 23:08.45 | 0.90 | $96.9\pm 0.5$ | 0.70 | $274.6\pm 0.7$ | 0.75 | $266.3\pm 1.1$ | $0.79\pm 0.35$
10b | 20.315 | 23:08.00 | 0.40 | $9.2\pm 0.6$ | 0.30 | $25.5\pm 0.7$ | 0.45 | $37.3\pm 0.9$ | $0.77\pm 0.35$
11b | 19.799 | 23:06.70 | 0.80 | $31.3\pm 0.3$ | 0.60 | $63.3\pm 0.7$ | 0.55 | $39.0\pm 0.3$ | $0.53\pm 0.24$
12b | 20.148 | 23:06.60 | 0.35 | $7.0\pm 0.4$ | 0.30 | $14.1\pm 0.7$ | 0.35 | $44.0\pm 1.1$ | $0.53\pm 0.25$
13b | 20.193 | 23:06.55 | 0.30 | $4.2\pm 0.4$ | 0.30 | $20.3\pm 0.7$ | 0.40 | $132.9\pm 1.6$ | $1.20\pm 0.54$
14b | 20.246 | 23:06.10 | 0.65 | $30.9\pm 1.0$ | 0.45 | $28.4\pm 1.1$ | 0.45 | $46.4\pm 3.0$ | $-0.06\pm 0.07$
15b | 20.133 | 23:06.05 | 0.35 | $6.9\pm 0.4$ | 0.25 | $12.7\pm 0.7$ | 0.30 | $59.0\pm 1.0$ | $0.46\pm 0.23$
16b | 20.174 | 23:05.80 | 0.40 | $11.8\pm 0.6$ | 0.30 | $45.9\pm 0.8$ | 0.50 | $267.1\pm 1.9$ | $1.03\pm 0.46$
17b | 19.693 | 23:05.65 | 0.65 | $15.9\pm 0.3$ | 0.40 | $20.2\pm 0.7$ | 0.40 | $13.6\pm 0.2$ | $0.18\pm 0.10$
18b | 20.072 | 23:05.10 | 0.45 | $12.7\pm 0.4$ | 0.30 | $22.0\pm 0.7$ | 0.40 | $45.5\pm 1.5$ | $0.42\pm 0.19$
19ab | 20.159 | 23:04.70 | 0.65 | $114.6\pm 0.8$ | 0.60 | $932.9\pm 1.3$ | 0.70 | $3194.7\pm 2.8$ | $1.59\pm 0.70$
20ab | 20.303 | 23:04.40 | 0.55 | $27.5\pm 1.0$ | … | … | 0.55 | $52.3\pm 2.7$ | …
21ab | 20.231 | 23:04.20 | 0.55 | $58.4\pm 0.8$ | 0.40 | $245.2\pm 0.9$ | 0.40 | $528.6\pm 1.7$ | $1.09\pm 0.48$
22a | 20.121 | 23:03.90 | 0.35 | $41.5\pm 0.3$ | 0.45 | $487.2\pm 0.9$ | 0.50 | $1353.7\pm 2.0$ | $1.87\pm 0.82$
23a | 20.174 | 23:03.50 | 0.45 | $47.1\pm 0.4$ | 0.40 | $225.3\pm 0.8$ | 0.40 | $517.8\pm 1.5$ | $1.19\pm 0.52$
24ab | 20.106 | 23:03.40 | 0.30 | $15.5\pm 0.3$ | … | … | 0.40 | $450.9\pm 1.6$ | …
25ab | 20.246 | 23:03.25 | 0.25 | $4.3\pm 0.3$ | 0.25 | $12.9\pm 0.7$ | 0.30 | $80.9\pm 1.1$ | $0.83\pm 0.38$
26b | 20.106 | 23:02.90 | 0.30 | $8.1\pm 0.2$ | 0.30 | $10.3\pm 0.7$ | 0.35 | $47.2\pm 1.2$ | $0.18\pm 0.12$
27ab | 20.277 | 23:02.85 | 0.45 | $24.0\pm 0.4$ | 0.45 | $198.2\pm 0.8$ | 0.50 | $387.9\pm 1.5$ | $1.60\pm 0.70$
28ab | 19.902 | 23:02.80 | 0.85 | $100.7\pm 0.3$ | 0.60 | $323.9\pm 0.8$ | 0.85 | $360.2\pm 0.6$ | $0.89\pm 0.39$
29b | 19.992 | 23:02.65 | 0.45 | $13.1\pm 0.2$ | 0.40 | $17.0\pm 0.7$ | 0.40 | $19.8\pm 0.7$ | $0.20\pm 0.10$
30b | 20.269 | 23:02.05 | 0.45 | $13.0\pm 0.3$ | 0.30 | $16.6\pm 0.7$ | 0.45 | $44.4\pm 1.0$ | $0.19\pm 0.10$
31ab | 20.125 | 23:02.00 | 0.95 | $56.3\pm 0.6$ | 0.55 | $62.1\pm 0.9$ | 0.55 | $73.3\pm 1.3$ | $0.07\pm 0.04$
32b | 19.898 | 23:01.90 | 0.35 | $11.3\pm 0.1$ | 0.30 | $18.5\pm 0.7$ | 0.45 | $40.5\pm 0.2$ | $0.37\pm 0.17$
33b | 19.860 | 23:01.35 | 0.45 | $10.4\pm 0.1$ | 0.30 | $19.8\pm 0.7$ | 0.45 | $24.1\pm 0.2$ | $0.49\pm 0.22$
34b | 20.125 | 22:60.00 | 0.70 | $22.7\pm 0.3$ | … | … | 0.75 | $41.6\pm 0.5$ | …
35∗ | 20.087 | 22:57.05 | 0.40 | $1.6\pm 0.1$ | … | … | 0.50 | $5.3\pm 0.2$ | …
36a | 19.598 | 22:56.25 | 0.75 | $30.9\pm 0.2$ | 0.70 | $38.7\pm 0.7$ | 0.55 | $26.7\pm 0.2$ | $0.17\pm 0.08$
37a | 20.186 | 22:56.25 | 0.40 | $3.1\pm 0.1$ | … | … | 0.40 | $3.9\pm 0.2$ | …
38a | 19.515 | 22:56.15 | 0.60 | $30.3\pm 0.2$ | 0.40 | $44.2\pm 0.7$ | 0.55 | $49.6\pm 0.2$ | $0.29\pm 0.13$
39∗ | 19.939 | 22:55.50 | 0.50 | $1.4\pm 0.1$ | … | … | 0.50 | $3.1\pm 0.2$ | …
40a | 19.526 | 22:55.30 | 0.65 | $36.6\pm 0.1$ | 0.40 | $43.0\pm 0.7$ | 0.55 | $45.6\pm 0.2$ | $0.12\pm 0.06$
41a | 18.629 | 22:54.30 | 1.25 | $73.9\pm 0.5$ | … | … | 1.15 | $84.5\pm 0.3$ | …
42∗ | 20.000 | 22:47.45 | 0.70 | $5.4\pm 0.1$ | … | … | 0.70 | $6.0\pm 0.2$ | …
43∗ | 19.886 | 22:47.40 | 0.60 | $3.1\pm 0.1$ | … | … | 0.60 | $5.4\pm 0.2$ | …
44b | 20.030 | 22:41.15 | 0.50 | $5.0\pm 0.1$ | 0.35 | $32.2\pm 0.7$ | 0.55 | $79.0\pm 0.2$ | $1.41\pm 0.62$
45b | 19.489 | 22:26.40 | 0.65 | $6.1\pm 0.1$ | … | … | 0.80 | $11.7\pm 0.2$ | …
46a | 19.803 | 22:20.70 | 1.10 | $79.1\pm 0.1$ | 0.65 | $227.5\pm 0.8$ | 0.50 | $141.0\pm 0.4$ | $0.80\pm 0.35$
47∗ | 19.417 | 22:19.75 | 0.60 | $3.0\pm 0.1$ | … | … | 0.60 | $7.3\pm 0.2$ | …
48a | 19.871 | 22:18.30 | 0.45 | $7.8\pm 0.1$ | 0.30 | $43.5\pm 0.7$ | 0.95 | $1401.0\pm 5.4$ | $1.30\pm 0.57$
49a | 19.898 | 22:17.00 | 0.65 | $28.2\pm 0.4$ | 0.45 | $184.2\pm 0.7$ | 0.50 | $290.9\pm 3.1$ | $1.42\pm 0.63$
50b | 18.102 | 22:06.85 | 0.50 | $6.4\pm 0.1$ | 0.25 | $8.9\pm 0.7$ | 0.50 | $9.7\pm 0.2$ | $0.25\pm 0.15$
51a | 19.996 | 22:04.65 | 1.20 | $162.6\pm 0.9$ | 0.80 | $411.3\pm 1.5$ | 1.00 | $271.4\pm 0.6$ | $0.70\pm 0.31$
52a | 17.337 | 22:03.60 | 0.75 | $16.5\pm 0.1$ | 0.45 | $29.8\pm 1.1$ | 0.65 | $19.8\pm 0.2$ | $0.45\pm 0.20$
53b | 23.049 | 21:55.55 | 0.85 | $23.6\pm 0.1$ | … | … | 0.90 | $30.3\pm 0.2$ | …
54a | 24.830 | 21:44.34 | 1.95 | $20.6\pm 0.1$ | … | … | 0.45 | $4.1\pm 0.2$ | …

${a}$: Associated with the Hii regions identified by Benson & Johnston (1984).
${b}$: Associated with the Hii regions identified by Gaume et al. (1995).
${*}$: Newly identified in this study.

Table 3: Derived parameters

# | $r_{\rm calc}$ | ${\rm EM}_{\rm calc}$ | $\log_{10}(\dot{N}_{\rm Ly}/{\rm s^{-1}})$ | $n_{\rm e}$ | $n_{\rm H_{2}}$ | $r_{i}$ | $t$
---|---|---|---|---|---|---|---
|  $\times 10^{-3}$ pc | $\times 10^{7}\,{\rm pc\,cm^{-6}}$ |  | $\times 10^{4}\,{\rm cm^{-3}}$ | $\times 10^{6}\,{\rm cm^{-3}}$ |  $\times 10^{-3}$ pc | $\times 10^{4}\,{\rm yrs}$
1 |  ∗26.73 |  ∗0.24 |  ∗46.17 |  ∗0.94 | … |  †∗0.106 |  †∗9.3
2 |  ∗57.40 |  ∗1.49 |  ∗47.64 |  ∗1.61 | 0.22 |  ∗0.324 |  ∗15.2
3 | 6.51 | 30.34 | 47.05 | 21.59 | … |  †0.207 |  †0.5
4 | 46.19 | 33.91 | 48.80 | 8.57 | 4.69 | 0.794 | 5.3
5 | 33.59 | 11.90 | 48.07 | 5.95 | … |  †0.453 |  †4.6
6 | 20.84 | 25.06 | 47.98 | 10.97 | 0.77 | 0.422 | 2.1
7 | 10.95 | 16.11 | 47.23 | 12.13 | … |  †0.237 |  †1.1
8 | 12.90 | 33.98 | 47.70 | 16.23 | … |  †0.339 |  †1.1
9 | 23.30 | 45.84 | 48.34 | 14.03 | 0.33 | 0.556 | 2.1
10 | 7.21 | 44.38 | 47.30 | 24.82 | 5.94 | 0.251 | 0.5
11 | 13.85 | 28.75 | 47.68 | 14.41 | … |  †0.337 |  †1.2
12 | 6.54 | 28.68 | 47.03 | 20.95 | 30.27 | 0.097 | 0.8
13 | 4.79 | 88.33 | 47.25 | 42.92 | 72.90 | 0.064 | 0.7
14 | 40.83 | 1.39 | 47.31 | 1.85 | 9.97 | 0.253 | 10.1
15 | 6.66 | 24.78 | 46.98 | 19.29 | 73.46 | 0.052 | 1.4
16 | 8.02 | 68.02 | 47.58 | 29.12 | 73.96 | 0.082 | 1.4
17 | 12.22 | 11.28 | 47.17 | 9.61 | … |  †0.227 |  †1.3
18 | 9.17 | 22.45 | 47.22 | 15.65 | 16.66 | 0.168 | 1.0
19 | 24.94 | 183.33 | 48.99 | 27.11 | 244.93 | 0.109 | 8.0
20 |  ∗20.36 |  ∗5.88 |  ∗47.33 |  ∗5.37 | 6.66 |  ∗0.257 |  ∗3.0
21 | 17.85 | 74.51 | 48.32 | 20.43 | 182.20 | 0.079 | 5.7
22 | 15.02 | 384.53 | 48.86 | 50.60 | 157.61 | 0.132 | 2.9
23 | 16.00 | 87.68 | 48.29 | 23.41 | 184.11 | 0.077 | 4.8
24 |  ∗8.63 |  ∗402.88 |  ∗48.39 |  ∗68.31 | 95.39 |  ∗0.129 |  ∗1.1
25 | 4.90 | 48.99 | 47.01 | 31.62 | 106.05 | 0.042 | 1.0
26 | 8.70 | 11.39 | 46.88 | 11.44 | 37.11 | 0.076 | 1.7
27 | 11.41 | 188.10 | 48.32 | 40.61 | 44.63 | 0.202 | 1.3
28 | 23.58 | 53.78 | 48.42 | 15.10 | 3.03 | 0.591 | 2.0
29 | 10.91 | 11.93 | 47.10 | 10.45 | 3.11 | 0.214 | 1.1
30 | 11.05 | 11.35 | 47.09 | 10.14 | 13.64 | 0.173 | 1.4
31 | 27.61 | 6.74 | 47.66 | 4.94 | 4.72 | 0.329 | 4.2
32 | 8.87 | 20.04 | 47.14 | 15.03 | 10.94 | 0.209 | 0.8
33 | 8.06 | 26.49 | 47.18 | 18.13 | 2.72 | 0.228 | 0.6
34 |  ∗26.73 |  ∗2.49 |  ∗47.20 |  ∗3.05 | 2.25 |  ∗0.231 |  ∗5.1
35 |  ∗13.66 |  ∗0.64 |  ∗46.02 |  ∗2.16 | 1.35 |  ∗0.094 |  ∗3.1
36 | 17.30 | 10.80 | 47.45 | 7.90 | … |  †0.282 |  †2.1
37 |  ∗13.66 |  ∗1.25 |  ∗46.31 |  ∗3.02 | 0.92 |  ∗0.117 |  ∗2.6
38 | 15.30 | 15.95 | 47.52 | 10.21 | 2.51 | 0.296 | 1.6
39 |  ∗18.18 |  ∗0.31 |  ∗45.95 |  ∗1.30 | 0.67 |  ∗0.089 |  ∗5.4
40 | 20.22 | 8.74 | 47.50 | 6.58 | 1.93 | 0.292 | 2.7
41 |  ∗49.30 |  ∗2.37 |  ∗47.71 |  ∗2.19 | 0.64 |  ∗0.342 |  ∗11.2
42 |  ∗26.73 |  ∗0.56 |  ∗46.54 |  ∗1.44 | 0.21 |  ∗0.140 |  ∗7.5
43 |  ∗22.51 |  ∗0.44 |  ∗46.30 |  ∗1.41 | 0.58 |  ∗0.116 |  ∗6.4
44 | 18.18 | 1.12 | 46.51 | 2.48 | 18.89 | 0.090 | 5.3
45 |  ∗24.63 |  ∗0.74 |  ∗46.60 |  ∗1.73 | 0.56 |  ∗0.146 |  ∗6.3
46 | 21.03 | 46.73 | 48.26 | 14.91 | … |  †0.522 |  †1.8
47 |  ∗22.51 |  ∗0.43 |  ∗46.28 |  ∗1.38 | 0.97 |  ∗0.115 |  ∗6.4
48 | 6.51 | 106.80 | 47.59 | 40.51 | 66.06 | 0.089 | 0.9
49 | 12.39 | 131.78 | 48.24 | 32.62 | 26.56 | 0.269 | 1.2
50 | 7.28 | 14.09 | 46.82 | 13.92 | 0.61 | 0.173 | 0.7
51 | 30.47 | 39.57 | 48.51 | 11.39 | … |  †0.633 |  †3.0
52 | 10.35 | 23.94 | 47.35 | 15.21 | … |  †0.261 |  †0.9
53 |  ∗32.96 |  ∗1.65 |  ∗47.20 |  ∗2.23 | 0.65 |  ∗0.232 |  ∗7.4
54 |  ∗77.55 |  ∗0.25 |  ∗47.12 |  ∗0.56 | … |  †∗0.218 |  †∗34.7
999$*$: Sources that are without detection in 22.4 GHz. Parameters are derived
from 6 GHz flux and the deconvolved core size. $\dagger$: No detection in 96
GHz. Parameters are derived assuming $n_{\rm H_{2}}=2\times 10^{7}\,{\rm
cm^{-3}}$

Table 4: Which cores in Ginsburg et al. (2018) are associates with outflows.

ID | Outflow | ID | Outflow | ID | Outflow | ID | Outflow | ID | Outflow | ID | Outflow
---|---|---|---|---|---|---|---|---|---|---|---
1 | Y | 47 | – | 97 | – | 143 | N | 202 | – | 258 | Y
2 | Y | 48 | – | 98 | – | 144 | N | 203 | – | 259 | Y
3 | Y | 49 | Y | 99 | – | 145 | – | 204 | – | 260 | Y
4 | Y | 50 | Y | 100 | – | 146 | – | 205 | – | 261 | N
5 | Y | 51 | Y | 102 | – | 147 | N | 206 | – | 262 | N
6 | N | 52 | – | 103 | – | 148 | – | 207 | Y | 263 | N
7 | Y | 53 | – | 104 | – | 149 | – | 208 | N | 266 | N
8 | Y | 54 | Y | 105 | – | 150 | – | 209 | N | 267 | N
9 | N | 55 | – | 106 | – | 153 | – | 210 | Y | 268 | N
10 | N | 56 | – | 107 | – | 154 | – | 211 | Y | 269 | N
11 | N | 57 | – | 108 | – | 155 | – | 212 | Y | 270 | N
12 | Y | 58 | – | 109 | – | 156 | – | 213 | – | 271 | N
13 | – | 59 | – | 110 | – | 157 | – | 214 | – | 101_X | –
14 | N | 60 | – | 111 | – | 158 | Y | 215 | – | 175_G | –
15 | – | 61 | – | 112 | – | 159 | Y | 216 | – | 177_B | –
16 | Y | 62 | – | 113 | – | 160 | Y | 217 | – | 180_E | –
17 | Y | 63 | – | 114 | – | 161 | Y | 218 | – | 181_D | –
18 | Y | 64 | – | 115 | – | 162 | Y | 219 | – | 182_Y | –
19 | Y | 65 | – | 116 | N | 163 | Y | 220 | – | 244_C | –
20 | Y | 66 | – | 117 | N | 164 | – | 221 | – | 265_H | N
21 | N | 67 | – | 118 | N | 165 | – | 222 | – | 171_K3 | –
22 | Y | 68 | – | 119 | – | 166 | – | 223 | – | 172_K2 | –
23 | Y | 69 | – | 120 | Y | 167 | – | 224 | – | 173_K1 | –
24 | Y | 70 | N | 121 | N | 168 | – | 226 | – | 174_f3 | –
25 | Y | 71 | N | 122 | N | 169 | – | 227 | – | 176_f1 | –
26 | – | 72 | N | 123 | N | 170 | – | 228 | – | 234_f4 | –
27 | – | 73 | N | 124 | N | 183 | N | 229 | – | 235_f2 | –
28 | – | 74 | N | 125 | N | 184 | N | 230 | – | 245_A2 | –
29 | – | 75 | N | 126 | N | 185 | N | 231 | – | 264_k4 | –
30 | – | 76 | N | 127 | N | 186 | N | 232 | – | 86_B9.96 | –
31 | – | 77 | N | 128 | N | 187 | N | 238 | – | 87_B9.99 | –
32 | – | 78 | N | 129 | N | 188 | N | 239 | – | 90_B9.89 | –
33 | – | 79 | N | 130 | – | 189 | N | 243 | – | 92_I10.52 | –
34 | Y | 80 | N | 131 | – | 190 | N | 246 | – | 96_Z10.24 | –
35 | Y | 81 | N | 132 | – | 191 | N | 247 | – | 151_B10.06 | –
36 | Y | 82 | N | 133 | – | 192 | N | 248 | N | 152_f10.32 | –
37 | N | 83 | N | 134 | – | 193 | N | 249 | N | 178_f10.37 | –
38 | Y | 84 | – | 135 | – | 194 | N | 250 | – | 179_f10.38 | –
39 | – | 85 | – | 136 | – | 195 | Y | 251 | – | 237_G10.44 | –
40 | – | 88 | – | 137 | – | 196 | Y | 252 | – | 241_f10.30 | –
41 | – | 89 | – | 138 | – | 197 | Y | 253 | – | 225_f10.33b | –
42 | – | 91 | – | 139 | – | 198 | Y | 254 | N | 233_f10.27b | –
43 | – | 93 | – | 140 | Y | 199 | – | 255 | N | 236_f10.303 | –
44 | – | 94 | – | 141 | Y | 200 | – | 256 | Y | 240_f10.44b | –
45 | Y | 95 | – | 142 | – | 201 | – | 257 | Y | 242_f10.318 | –
46 | Y |  |  |  |  |  |  |  |  |  |

101010The IDs of the dust cores are as in Table 3 in Ginsburg et al. (2018).
Cores that are not covered by the SiO image (Fig. 15) are marked “–”. Cores
that are covered by the SiO image and associated or not associated with
outflows are flagged Y or N, respectively.

##  Appendix C SiO Images

Figure 14: Peak intensity map of SiO (5-4) line. The zoomed-in plots are the
regions in Sgr B2(DS) with abundant outflows. The resolution of the image is
0.35″$\times$ 0.24″, with P.A. of $-80^{\circ}$. The dust cores identified by
Ginsburg et al. (2018) are shown as green circles, whose size is not scaled to
the size of the cores.

Figure 15: Moment 1 map of SiO (5-4) line. The zoomed-in plots are the regions
in Sgr B2(DS) with abundant outflows. The resolution of the image is
0.35″$\times$ 0.24″, with P.A. of $-80^{\circ}$. The dust cores identified by
Ginsburg et al. (2018) are shown as green circles, whose size is not scaled to
the size of the cores.