Sgr B2(N): A bipolar outflow and rotating hot core revealed by ALMA

Abstract

We present the results of SiO (2 − 1) and SO 2 (12 4,8 − 13 3,11) line observations of Sgr B2(N) made with the Atacama Large Millimeter/submillimeter Array (ALMA) at an angular resolution of ∼2 ′′. Our analysis of the SiO and SO 2 line emission reveals a bipolar molecular outflow in an east-west direction whose driving source is located at K2. In addition, SO 2 line core shows a north-south velocity gradient most probably indicating a hot core of molecular gas rotating around K2. Fractional abundances of SO 2 and SiO (X(SO 2) and X(SiO), respectively) in the outflowing molecular gas are derived from comparisons with the C 18 O emission. Assuming an excitation temperature of 100 ± 50 K, we calculate X(SO 2) = 2.3 +2.6 −0.4 ×10 −8 and X(SiO) = 1.2 +0.1 −0.1 ×10 −9. The outflow from Sgr B2(N) K2 is characterized as a young (5×10 3 yr) and massive (∼2000 M ⊙), but moderately collimated (∼60 •) outflow. We also report a possible detection of the SiO (v = 2, J = 2 − 1) maser emission from the position of K2. If confirmed, it would make Sgr B2(N) the 4 th star forming region associated with SiO masers. Subject headings: ISM: kinematics and dynamics — ISM: molecules — ISM: individual (Sgr B2) — ISM: outflows — stars: massive — stars: formation

Full text

Sgr B2(N): A BIPOLAR OUTFLOW AND ROTATING HOT CORE REVEALED BY ALMA
Aya E. Higuchi
1
, Tetsuo Hasegawa
2
, Kazuya Saigo
3
, Patricio Sanhueza
2
, and James O. Chibueze
4
1
College of Science, Ibaraki University, 2-1-1 Bunkyo, Mito, 310-8512, Japan; aya.higuchi.sci@vc.ibaraki.ac.jp
2
National Astronomical Observatory of Japan 2-21-1 Osawa, Mitaka, Tokyo, 181-8588, Japan
3
Department of Physical Science, Graduate School of Science, Osaka Prefecture University, 1-1 Gakuen-cho, Naka-ku, Sakai, Osaka 599-8531, Japan
4
Department of Physics and Astronomy, Faculty of Physical Sciences, University of Nigeria, Carver Building, 1 University Road, Nsukka, Nigeria
Received 2015 July 29; accepted 2015 November 9; published 2015 December 15
ABSTRACT
We present the results of SiO (2−1)and SO
2
(12 13
4,8 3,11
–)line observations of Sgr B2(N)made with the Atacama
Large Millimeter/submillimeter Array at an angular resolution of ∼2″. Our analysis of the SiO and SO
2
line
emission reveals a bipolar molecular outflow in an east–west direction whose driving source is located at K2. In
addition, the SO
2
line core shows a north–south velocity gradient most probably indicating a hot core of molecular
gas rotating around K2. Fractional abundances of SO
2
and SiO (X(SO
2
)and X(SiO), respectively)in the outflowing
molecular gas are derived from comparisons with the C
18
O emission. Assuming an excitation temperature of
100±50 K, we calculate X(SO
2
)=
2
.3 0.4
2.
6
-
+×10
−8
and X(SiO)=1.2 0.1
0.1
-
+×10
−9
. The outflow from Sgr B2(N)K2
is characterized as a young (5×10
3
yr)and massive (∼2000 M)but moderately collimated (∼60°)outflow. We
also report a possible detection of the SiO (vJ2, 2 1–==
)maser emission from the position of K2. If confirmed,
it would make Sgr B2(N)the fourth star-forming region associated with SiO masers.
Key words: ISM: individual objects (Sgr B2)–ISM: jets and outflows –ISM: kinematics and dynamics –
ISM: molecules –stars: formation –stars: massive
1. INTRODUCTION
High-mass young stellar objects are usually deeply
embedded in their parental dense and massive molecular
clumps (size∼1 pc, mass∼100–1000 M
,
density ∼10
45-
cm
3-
; Lada & Lada 2003; Ridge et al. 2003; Higuchi et al.
2009,2010,2013; Lada 2010), obscuring their early formative
stages. Their formation timescales of ∼10
5
yrare short, and
they form in distant clusters (e.g., Galactic center)and
associations (e.g., Zinnecker & Yorke 2007). All these factors
limit our understanding of their formation processes. High
angular resolution observations are indispensable in the efforts
to unveil the mystery of high-mass star formation. The
Atacama Large Millimeter/submillimeter Array (ALMA)
provides the high sensitivity, angular resolution, and dynamic
range to improve our understanding of the formation processes
of high-mass stars and their parental clumps (Sánchez-Monge
et al. 2013,2014; Belloche et al. 2014; Guzmán et al. 2014;
Higuchi et al. 2014,2015; Johnston et al. 2015; Zapata
et al. 2015).
Sagittarius B2 (Sgr B2)is a complex of H II regions and giant
molecular clouds located near the Galactic center. It is known as
one of the most luminous massive star-forming regions in our
Galaxy, with a total bolometric luminosity of ∼10
7
M(Lis &
Goldsmith 1989,1990,1991). Tens of compact/ultracompact H
II regions and H
2
O, OH, and H
2
CO maser clusters are located
at the three centers of star formation aligned in the north–
south direction, i.e., Sgr B2(N), Sgr B2(M), and Sgr B2(S)
(Genzel et al. 1976; Benson & Johnston 1984; Gardner
et al. 1986; Gaume & Claussen 1990). Sgr B2(N)is associated
with the H II region K, which is further resolved into
subcomponents K1–K6 (Gaume & Claussen 1990; Gaume
et al. 1995; de Pree et al. 1996). Of these, K1, K2, and K3 have
drawn particular attention because they are very compact
(03–0 6)(Gaume et al. 1995)and are spatially coincident with
the H
2
O masers (Kobayashi et al. 1989).
A dense core of hot molecular gas was initially found in the
K1–K3 region by interferometric observations of NH
3
and
HC
3
N line emission (Vogel et al. 1987; Gaume & Claussen
1990; Lis et al. 1993). The core is also detected in dust
continuum emission at millimeter wavelengths. The dust
emission peaks at the position of K2 with the peak H
2
column
density of
N
H2∼(0.4–8)×10
25
cm
−2
(Carlstrom &
Vogel 1989; Lis et al. 1993; Kuan et al. 1996). The inferred
mass of the core ranges from 104to 105M
,
for a region with a
diameter of 0.2–0.4 pc. Interferometric and single-dish follow-
up observations in several molecular lines have revealed that
the core is particularly rich in large saturated molecules,
indicating hot core chemistry with fresh material evaporating
from dust mantles (Goldsmith et al. 1987; Kuan &
Snyder 1994; Kuan et al. 1996; Liu & Snyder 1999; Hollis
et al. 2003; Belloche et al. 2008,2013,2014; Qin et al. 2011).
For the kinematic structure of the molecular gas in this
region, two possible interpretations have been presented. Vogel
et al. (1987)were the first to note the NW–SE velocity gradient
in their NH
3
observations. Lis et al. (1993)confirmed the
velocity structure in their HC
3
N(25–24)observations at ∼4″
resolutionand suggested that it arises from a very energetic
bipolar outflow associated with a source embedded in the dense
core. On the other hand, Hollis et al. (2003)argued, based on
their 1 5 resolution observation of the CH
3
CH
2
CN (
5
4
1,5 1,4
–)
line, that this velocity gradient indicates the rotation of a large
edge-on disk that extends in the east–west direction. Our
understanding of the kinematics remains unsettled between
these two contradicting pictures.
In this paper, we analyze the ALMA archival data taken at
∼2″resolution and present images of Sgr B2(N)in SiO (2–1)
(hereafter SiO)and SO
2
(12 13
4,8 3,11
–)(hereafter SO
2
)line
emission. At the adopted distance of 7.8 kpc (Reid et al. 2009),
2″corresponds to 0.076 pc. The SO
2
molecule is a good tracer
of dense and hot gas typical of hot cores (e.g., Charnley 1997;
van der Tak et al. 2003; Jiménez-Serra et al. 2007; Leurini et al.
2007), while SiO is an excellent tracer of molecular outflows
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 doi:10.1088/0004-637X/815/2/106
© 2015. The American Astronomical Society. All rights reserved.
1

and shocks (e.g., Downes et al. 1982; Bachiller et al. 1991;
Zapata et al. 2009a; Leurini et al. 2014). The goal is to establish
a clearer kinematic picture of the Sgr B2(N)coreand to enable
further explorations of the processes associated with the
ongoing high-mass stars formation. We also report the possible
discovery of the SiO (v=2, J=2–1)maser emission from
Sgr B2(N).
2. OBSERVATIONS
Sgr B2(N)was observed with ALMA (Hills et al. 2010)
during Early Science Cycle 0 (Belloche et al. 2014). The
observations were done in one execution with 2112m
antennas and three executions with 2612m antennas at an
angular resolution of ∼2The maximum baseline length
achieved during the observations was 440 m. The full
frequency range between 84 and 111 GHz was covered in the
spectral scan mode with four overlapping basebands. The
HPBW of the primary beam varies between 74″at 84 GHz and
56″at 111 GHz. The channel spacing was 244 kHz, and
smoothed spectra were 488 kHz (∼1.6 km s
−1
).
The ALMA calibration includes simultaneous observations
of the atmospheric 183 GHz water line with water vapor
radiometers. The measured water columns in the antenna beam
were later used to reduce the atmospheric phase fluctuations.
Amplitude calibration was done using Neptune. The quasars
B1730-130 and J1700-261 were used to calibrate the bandpass
and the complex gain fluctuations, respectively. Data reduction
was performed using the Common Astronomy Software
Applications package (CASA; http://casa.nrao.edu). Natural
weighting was used for imaging. The sensitivity obtained
ranges from 3 to 7 mJy beam
−1
per channel of 1.6 km s
−1
(see
details in Table 1).
In this paper, we focus on the SO
2
and SiO lines. Among
several transitions of SO
2
detected between 84 and 111 GHz in
the ALMA observations, we select the
124,8–133,11 transition
for the majority of our analysis because(1)its upper-state
energy is high (
E
110.6
u=K)and it traces the hotter and
denser region in the immediate vicinity of the young star, and
(2)this transition is relatively less contaminated by the other
coincident lines so that we can get a clearer kinematic picture
of the core. SiO is an excellent tracer of shocks and molecular
outflows (e.g., Downes et al. 1982; Martín-Pintado et al. 1997;
Zapata et al. 2009a). We also investigate the lines of SiO in the
v=1, 2, 3 vibrationally excited states. Parameters of the data
used are summarized in Table 1.
3. RESULTS AND DISCUSSION
3.1. Line Profiles
Figure 1shows the spectra of the SO
2
and SiO lines
integrated within a 20″×20″region centered on the K2
position: (,
J2000
aJ200
0
d
)=[
17
h
4
7m19. 88,
s
−28°22′18 4]
(Gaume et al. 1995). Many spectral lines other than SO
2
and
SiO are detected in the ALMA data toward Sgr B2(N). We also
label the molecular lines near the SO
2
and SiO lines identified
by Belloche et al. (2013)using the IRAM 30 m telescope.
Although these molecular lines partially overlap with the SiO
and SO
2
lines, we can separate them because their spatial
distributions are different.
The integrated SO
2
spectrum in Figure 1(a)has an emission
peak of 4.0Jy at 63 km s
−1
, which we adopt for the systemic
velocity (VSYS)of Sgr B2(N).Wedefine V
D
as V
D
ºVLSR–
VSYS (VSYS =63 km s
−1
). By visual inspection of the SO
2
spectrum, we find a line core from
V
LSR of 56 to 70 km s
−1
(V
D
=−7to+7kms
−1
), a blueshifted wing from
V
35
LSR »
to 56 km s
−1
(V28
D
»- to −7kms
−1
), and a redshifted wing
from
V
70
LSR »to 110 km s
−1
(V7
D
»+ to +47 km s
−1
).
The integrated SiO spectrum in Figure 1(b)shows a deep
absorption from
V
48
LSR »to 92 km s
−1
(V15
D
»- to
+29 km s
−1
). This is consistent with the observations made
by Liu et al. (1998)using the Berkeley–Illinois–Maryland
Association (BIMA)array. In addition, the present SiO
spectrum exhibits wing emission at blueshifted velocities of
V
30
LSR »–48 km s
−1
(V33
D
»- to −15 km s
−1
)and red-
shifted velocities of
V
9
2
LSR »–110 km s
−1
(V29
D
»+ –
47 km s
−1
).
3.2. Kinematic Structure of SO
2
and SiO Emission
3.2.1. SO
2
Line Emission
Figure 2shows the channel maps of SO
2
emission in the
velocity range of
V
LSR =47–85 km s
−1
(V
D
=−16 to
+22 km s
−1
). The positions are shown in offsets from the
K2. At the velocities of the blueshifted wing shown in Figure 2
(
V
LSR =47–56 km s
−1
,V
D
=−16 to −7kms
−1
)the SO
2
emission extends E–SE of K2, while at the velocities of the
redshifted wing in Figure 2(
V
LSR =70–85 km s
−1
,V
D
=+7
to +22 km s
−1
)the emission extends W–NW of K2. This
makes a clear E–W velocity gradient. However, within
the “line core”(
V
LSR =56–70 km s
−1
; also see Figure 1),
the velocity gradient changes its direction. At
V
LSR =
57.8–61 km s
−1
, the peak of the SO
2
emission islocated
S–SE of K2, while at
V
LSR =65.8–69 km s
−1
the peak is NE–N
Table 1
Parameters for the ALMA Observations of Sgr B2(N)
Molecule Transition νSynthesized Beam Velocity Resolution rms Noise Level
(GHz)(km s
−1
)(mJy beam
−1
)
SO
2
1
24,8–133,11 107.843 1 9 ×1 4 1.6 6
SO
2
2
94,26–285,23 99.393 1 9 ×1 4 1.6 6
SO
2
3
25,27–316,2
6
84.321 2 2 ×1 6 1.6 7
SiO v=0, J=2–1 86.847 2 5×1 8 1.6 (spectrum)and 4.8 (channel maps)7(channel maps)
SiO v=1, J=2–1 86.243 2 5×1 8 1.6 3
SiO v=2, J=2–1 85.640 2 5 ×1 8 1.6 3
SiO v=3, J=2–1 85.038 2 5 ×1 8 1.6 3
Note. The rms noise level is derived in the emission-free area defined by inspecting the channel maps.
2
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

of K2. This “line core”component may predominantly
originate from the central compact core and is likely tracing
a different kinematic structure (e.g., Guzmán et al. 2014;
Sánchez-Monge et al. 2014). A secondary peak is seen ∼5″
north of K2 at
V
LSR∼70 km s
−1
. This is identified with one of
the compact “quasi-thermal cores,”source h, observed in
CH
3
OH by Mehringer & Menten (1997).
The change of the direction of the velocity gradient
described above is clearly shown by the moment maps in
Figure 3. The first-moment map (intensity-weighted velocity)
made from the entire velocity range of
V
LSR =33–109 km s
−1
(V
D
=−30 to +46 km s
−1
)shows the E–W velocity gradient
(Figure 3(a)). In contrast, the first-moment map made from
channels near the systemic velocity (
V
LSR =59–65 km s
−1
)
shows the S–N velocity gradient (Figure 3(b)).
Figure 3(c)shows the positional shift of the SO
2
emission
peak with velocity. As we noted in Figure 2, the peaks at the
wing velocities (V
∣∣
D>7kms
−1
)tracethe E–W velocity
gradient, while the peak positions at the “line core”velocities
(V
∣∣
D<7kms
−1
)clearly showthe N–S velocity gradient
perpendicular to the former.
3.2.2. SiO Line Emission
Figure 4shows the velocity channel maps of the SiO wing at
emission seen in Figure 1(b), i.e., the blueshifted wing at
V
LSR =32–51 km s
−1
(V
D
=−31 to −12 km s
−1
)and the
Figure 1. Integrated spectra around the (a)SO
2
and (b)SiO emission lines in Sgr B2(N). The spectra were integrated over a 20″×20″region, centered on the
continuum source K2. VLSR of 63 km s
−1
is the systemic velocity of Sgr B2(N).
3
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

redshifted wing at
V
LSR =90–114 km s
−1
(V
D
=+27 to
+51 km s
−1
).
The SiO emission at the blueshifted wing extends E–SE of
K2, while that at the redshifted wing extends W–NW of K2.
This is similar to what we see for the SO
2
emission at the wing
velocitiesand suggests that SiO emission and SO
2
emission
share the same kinematics.
3.3. Bipolar Outflow and Rotating Hot Core
As we described in Section 1, the understanding of the
kinematic structure of the Sgr B2(N)core is unsettled. Two
mutually contradicting interpretations have been put forward so
far:(a)an E–W bipolar jet and a disk rotating in the N–S
direction (e.g., Lis et al. 1993), and (b)a large disk rotating in
Figure 2. Velocity channel maps of the SO
2
emission in white contours and color image for Sgr B2(N). The velocity intervals are 1.6 km s
−1
. Contours start from
−5σ,5σ,10σ,to30σ, increasing in intervals of 10σlevels, and then they continue in steps of 30σup to the 150σlevel (1σ=6 mJy beam
−1
). Negative contour levels
are shown with dashed lines. The red crosses mark the position of K2 (Gaume et al. 1995).
4
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

the E–W direction (Hollis et al. 2003). What do the new
ALMA observations tell us?
The SO
2
and SiO spectra show the wing emission at high
velocities, which suggests the presence of an outflow with
velocities far exceeding 20 km s
−1
(V28
D
»- to +47 km s
−1
for SO
2
and V33
D
»- to +47 km s
−1
for SiO). Figure 5(a)
shows the spatial distribution of the SiO wing emission that we
identify as the bipolar outflow. The velocity ranges were
selected to separate distinct feature of outflow lobes. The
bipolar outflow lobes are aligned in the E–W direction
(P.A. =120°±10°)and are symmetrically displaced about
the peak of the integrated SO
2
emission near K2. Figure 5(b)
shows the zoom-up of the central region. Blue and red crosses
mark the position of the H
2
O maser spots reported by McGrath
et al. (2004)blue- and redshifted with respect to the systemic
velocity
V
SYS =63 km s
−1
, respectively. The position and
velocity of the masers are consistent with the innermost part of
the blue- and redshifted outflowing gas traced by SO
2
and SiO.
The white cross marks the position of the possible SiO maser
source presented in Section 3.5.
Figure 6(a)shows the observed position–velocity (PV)
diagram of the SO
2
emission along the major axis of the
bipolar outflow (P.A. =120°). The PV diagram along this
axis shows blueshifted and redshifted emission extending
to V28
D
»- to +47 km s
−1
. Figure 6(b)shows the observed
PV diagrams of the SO
2
emission in the direction perpendicular
to the axis of the outflow (P.A. =30°). Between VLSR =59
and65 km s
−1
, the peak velocity changes almost linearly as a
function of the position with a constant velocity gradient.
Figure 6(c)shows the PV diagram of the SiO emission along
the major axis of the blue- and redshifted components
(P.A. =120°). The blue- and redshifted SiO emission is seen
extending up to V33
D
»- to +47 km s
−1
. This shows that
the SiO and SO
2
lines share the same kinematics except for
the severe foreground absorption in SiO at low velocities
(V
D
=−15 to +29 km s
−1
).
From the results presented above, we summarize the
evidences for the existence of the bipolar outflow as
follows:(1)presence of high-velocity gas (wing components)
in the spectra of shock tracers, (2)the bipolar structure is highly
symmetric and well traced by the SiO line, which is a good
tracer of outflows, (3)H
2
O masers are associated with the
bipolar structure, and (4)the high-velocity motion cannot be
gravitationally bound by the central core. In fact, were the high-
velocity emission a result of a gravitationally bound circular
motion, the mass required inside its orbit would be estimated
Figure 3. (a)First-moment map of SO
2
emission made using all velocity ranges (33–109 km s
−1
).(b)First-moment map of SO
2
emission made using velocity ranges
of 59–65 km s
−1
.(c)Zoom-inof the SO
2
peak position of panel(b). The red crosses mark the position of K2 (Gaume et al. 1995).
5
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

by the dynamical mass M;
dyn
Mr
G,1
dyn
23
()
w
=
where ωis an angular velocityand ris a rotational radius. From
the observed angular velocity of 1.6 ×10
−11
s
−1
and a 0.08 pc
radius, we derive a dynamical mass of ∼3×10
4
Mneeded
to gravitationally bind the motion, which should be seen as
a lower limit because we do not correct for the inclination of
the motion. This value is higher than the mass of the core
derived from the 1.3 and 1.1 mm continuum emission with
larger core radius (9×10
3
Mby Lis et al. 1993;1×10
4
M
by Qin et al. 2008;(4–20)×10
3
Mby Liu & Snyder
1999)and shows that the gas responsible for the high-velocity
wing emission cannot be gravitationally bound by the
central core.
Meanwhile, we consider that the N–S velocity gradient
observed in SO
2
at P.A. =
3
0,
◦i.e., orthogonal to the bipolar
Figure 4. Velocity channel maps of the SiO emission in white contours and color image for Sgr B2(N). Contours start from −5σ,5σ,10σ,to60σ, increasing in
intervals of 10σlevels (1σ=7 mJy beam
−1
). Negative contour levels are shown with dashed lines. The velocity intervals are 4.8 km s
−1
. The red crosses mark the
position of K2.
6
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

outflow, is a signature of the rotating motion of the core
(see Figure 6(b)). It has been reported that SO
2
traces rotating
cores in other regions of massive star formation (e.g., Beltrán
et al. 2014; Guzmán et al. 2014). From Figure 6(b),we
determine a velocity gradient of 88 km s
−1
pc
−1
or an angular
velocity 2.8 10 s
12 1
w
=´
--
at a radius
r
0.028=pc (or
5800 AU). Using Equation (1), we estimate that a dynamical
mass of 42Misin 2
()
is needed within this radius to
gravitationally bind the rotating motion, where iis the
inclination of the rotation axis with respect to the line of sight.
Our conclusion on the configuration of the bipolar outflow
and the rotating core supports the original interpretation by Lis
et al. (1993). They noted the N–S linear velocity gradient in the
core of ∼2kms
−1
arcsec
−1
(1.7 10 1
2
w
=´
-s
−1
at the
distance of 7.8 kpc), which is 40% smaller than our estimate
above. This small discrepancy may be because of their larger
synthesized beam (45×3 7 FWHM)and the fact that they
measured the gradient in the N–S direction while the gradient is
steepest at P.A. =
3
0.
◦
The SO
2
emission in Figure 6(b)exhibits a constant velocity
gradient rather than the pattern of the Keplerian rotation that is
characterized by an increase of the rotation velocity at closer
distance to the central star. To further examine this point, we
analyzed the SO
2
2
94,26–285,23 (99.3925 GHz)and
3
25,27–316,26
(84.3209 GHz)emissions that arise from energy levels at
E
44
1
u=and 549 K from the ground state, respectively, much
higher than the
E
110.6
u=K for the 124,8–133,11 transition we
have discussed so far. Both lines are free from contamination
by other molecular linesand are detected at the peak flux
density of 0.44Jybeam
−1
(294,26 –285,23)and 0.28Jybeam
−1
(325,27–316,26), with the distributions quite similar to each
other. Figure 6(d)shows the PV diagram of the SO
2
2
94,26–
285,23 emission at P.A. =30°(color)in comparison with that of
the SO
2
124,8–133,11 emission (contours). Although the
2
94,26–
285,23 emission arises from much hotter and denser molecular
gas, the velocity gradient is similar to that of the 124,8–133,11
emission. This situation is quite different from the case of a
disk in Keplerian rotation around a massive star, for which we
should see a steeper velocity gradient for lines arising from the
hotter and denser region closer to the star. We conclude that the
“line core”emission of SO
2
in Sgr B2(N), when observed at a
∼2″(0.076 pc or 15,600 AU)resolution, arises from a rotating
ring-like structure with a radius of ∼6000 AU.
3.4. Physical Parameters of the Bipolar Outflow
3.4.1. Fractional Abundances of SO
2
and SiO
In order to estimate the physical parameters of the outflow,
we need to know the fractional abundances of SO
2
and SiO,
defined as X(SO
2
)=[SO
2
]/[H
2
]and X(SiO)=[SiO]/[H
2
],
which are known to vary significantly from one object to
another (e.g., Martin-Pintado et al. 1992; Charnley 1997; van
der Tak et al. 2003; Sanhueza et al. 2012,2013). We estimate
the SO
2
and SiO fractional abundances by a comparison with
the Sgr B2(N)outflow detected in C
18
O, which is one of the
molecular probes with stable fractional abundances. Unfortu-
nately, an analysis of the C
18
O emission included in the present
observations (Hasegawa et al. 2015, in preparation)shows that
overlaps with the emission of HNCO (v
5
=1),CH
3
CN
(v
4
=1), and C
2
H
5
OCHO (e.g., Belloche et al. 2013)lines
make it difficult to analyze the high-velocity C
18
O wing
emission over the full velocity range of the outflow. Instead,
we choose a less contaminated velocity channel at 46.6 km s
−1
and compare with SO
2
and SiO. Even at this velocity,
point-like CH
3
CN(v
4
=1)emission at K2 contaminates the
C
18
O image. We avoid the K2 position and average C
18
O, SO
2
,
and SiO emission over a 6″×6″box centered at
(,
J2000
aJ200
0
d
)=[
17
h
4
7m
2
0 . 22,
s
−28°22′19 7], which is
∼4″away from the position of K2.
For derivation of the SO
2
and SiO abundances, we assume
that SO
2
, SiO, and C
18
O lines are optically thin and in a
local thermodynamical equilibrium (LTE, i.e., they are excited
to a common excitation temperature TTCO
ex ex 18
()==
TTSO SiO
ex 2 ex
() ()=). In addition, we assume that their
relative abundances are uniform in the PV space defined by
the outflow. For T,
ex we have adopted a range of possible
temperatures (50, 100, and 150 K)to compare the results.
Figure 5. (a)SO
2
integrated intensity map (grayscale image and black
contours)and the integrated intensity map of the SiO outflows (blue and red
contours)of Sgr B2(N). The blue color represents blueshifted gas, while the red
color represents the redshifted gas. The black contours range from 10% to 90%
of the peak emission, in steps of
1
0%
.
The blue color represents blueshifted gas
(VLSR =29–50 km s
−1
), while the red color represents the redshifted gas
(VLSR =91–111 km s
−1
). The contours for integrated intensity maps, with
intervals of 5σ, start from the 5σlevel (1σ=0.11Jybeam
−1
km s
−1
for both
blue- and redshifted gas). The grayscale bar shows the flux density of the SO
2
emission. The yellow cross marks the K2 position. (b)Zoom-in ofthe SO
2
integrated intensity map and the SiO outflow. The blue and red crosses mark
the position of the blue- and redshifted water maser spots reported by McGrath
et al. (2004), respectively. The white cross marks the position of the possible
SiO maser (see Section 3.5).
7
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

Molecular column densities can be calculated (Liu et al.
1998)from
N
QT
SSdv2.04 10
exp
,2
E
T
ab
20 ex
32
u
ex
()
()
()
()
ò
qqnm
=´ ´n
where ab
q
q´are the FWHM major and minor axes of the
synthesized beam in units of arcseconds,
Q
Tex
()
is the partition
function, E
u
is the upper energy level in K, νis the rest
frequency of the transition in GHz, and
S
2
mis the product of
the intrinsic line strength and the squared dipole momentum in
D
2
.
S
nis the measured intensity in Jybeam
−1
, and Tex is the
excitation temperature in K. Sdv
ònis the measured integrated
line emission in Jybeam
−1
km s
−1
.
For the SO
2
line, we use Q(150 K)=2091, Q(100 K)=
1140, andQ(50 K)=404 from Pickett et al. (1998),
E
u
=110.6 K, and
S
2
m=8.27D
2
(Belloche et al. 2013).
Similarly for the SiO line, we use Q(150 K)=144,
Figure 6. PV diagrams along the outflow in SO
2
and SiO emission, P.A. =120°(panels (a)and (c)), and perpendicular to the outflow in SO
2
emission, P.A. =30°
(panels (b)and (d)).(a)Contour levels [5, 10, 30, 60, 90, 120, 150, 180, 210, 240, 270, and 300]×1σ(1σ=3 mJy beam
−1
).(b)Contour levels [5, 10, 30, 60, 90,
120, and 150]×σ(1σ=5 mJy beam
−1
).(c)contour levels [5, 10, 15, 20, 25, 30, 35, and 40]×σ(1σ=10 mJy beam
−1
).(d)SO
2
(294,2
6
–285,23)(color)and
SO
2
(124,8 –133,11)(contours), contours levels [5, 10, 30, 60, 90, 120, and 150]×σ(1σ=5 mJy beam
−1
). Other molecular lines (e.g., C
2
H
5
OH, n-C
3
H
7
CN)are also
displayed.
8
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

Q(100 K)=96, Q(50 K)=48 (Pickett et al. 1998),E
u
=
6.25 K, and
S
2
m=19.2D
2
(Liu et al. 1998; Fernández-López
et al. 2013). The calculated column densities of SO
2
and SiO
per unit velocity width at
V
46.6
LSR =km s
−1
are shown in
Table 2with the assumed excitation temperatures. For the
fractional abundances, we derive the H
2
column densities per
unit velocity width from the C
18
O column densities per unit
velocity width at
V
46.6
LSR =km s
−1
and the adopted X(C
18
O)
of 110
7
´-for Sgr B2 (Lis & Goldsmith 1989).
Table 2shows the SO
2
and SiO fractional abundances
derived for Tex =50, 100, and 150K. We note that the
fractional abundances vary only mildly with the assumed LTE
temperature, particularly for SiO. For the following discussion
of the physical parameters of the outflow, we adopt the
values derived for T100
ex =K with an uncertainty range
estimated from the cases of Tex =50 and 150 K, i.e.,
X(SO
2
)=
2
.3 0.4
2.
6
-
+×10
−8
and X(SiO)=1.2 0.1
0.1
-
+×10
−9
.
These values are consistent with the previous estimates for
other high-mass hot cores. Van der Tak et al. (2003)
andEsplugues et al. (2013)derived molecular abundances
of SO
2
in high-mass star-forming regions ranging from
X(SO
2
)=10
−6
to 10
−8
. Gusdorf et al. (2008), Tercero et al.
(2011), Sanhueza et al. (2013), and Leurini et al. (2014)found
fractional abundances of SiO in the high-mass protostellar
outflows ranging from X(SiO)=10
−7
to 10
−9
. We should keep
in mind that there is an uncertainty in our determination of the
fractional abundances due to the assumption of uniform
excitation and chemistry over space and velocity, which is
the best allowed by the present data, although it is obviously a
bold simplification.
3.4.2. Physical Parameters
Using the derived fractional abundances, we estimate the
mass of the outflow material for individual velocity channels of
the the SO
2
and SiO images. Figure 7shows the plots of (a)
masses, (b)momenta, and (c)energies in the outflow per unit
velocity width as a function of velocity offset from the systemic
velocity (V
∣∣
D=VLSR
∣
–
V
SYS
∣
). The mass plot in Figure 7(a)
suggests an approximate power law in the form of
d
Mv()dvµV,
∣
∣Dgwith a power-law index γ∼−1. This
slope is quite shallow compared with cases of outflows from
low-mass protostars (4
g
~- ; Shepherd et al. 1998). Richer
et al. (2000)and Arce et al. (2007)noted a tendency for the
slope to steepen (from 1
g
~- to −10)with the outflow age.
The shallow slope of the Sgr B2(N)outflow is consistent with
its youth (t510
dyn 3
~´ yr;see Table 3).
Figures 7(b)and (c)show that the momentum is distributed
rather evenly over the velocity range of the outflow, and that
the majority of the kinetic energy is carried at larger V.
∣
∣DIt
would be quite interesting to see how these plots compare with
similar plots for other regions of high-mass star formation seen
at high spatial resolution.
Table 3summarizes the physical properties of the molecular
gas in the outflow. The momentum is given by ΣM
i
V
i
∣∣
Dand
the energy by (1/2)ΣM
i
V,
i2
∣
∣Dwhere M
i
is the outflow mass
in the velocity channel iand V
i
D
is its velocity offset relative
to
V
.
LSR For the derivation of the outflow parameters, no
correction for inclination angle was applied. We integrate only
the velocity ranges of V
∣∣
D>7kms
−1
for deriving the
Table 2
Column Densities and Fractional Abundances in the Outflowing Gas
a
Parameter Tex =50 K Tex =100 K Tex =150 K
N
dv
SO2(cm
−2
/km s
−1
)
b
8.4 ×10
14
7.8 ×10
14
9.9 ×10
14
N
dv
SiO (cm
−2
/km s
−1
)
b
2.2 ×10
13
4.2 ×10
13
6.2 ×10
13
N
dv
CO
18 (cm
−2
/km s
−1
)
c
1.7 ×10
15
3.5 ×10
15
5.2 ×10
15
N
dv
H2(cm
−2
/km s
−1
)
d
1.7 ×10
22
3.5 ×10
22
5.2 ×10
22
X(SO
2
)4.9 ×10
−8
2.3 ×10
−8
1.9 ×10
−8
X(SiO)1.3 ×10
−9
1.2 ×10
−9
1.2 ×10
−9
Notes.
a
Measured at V46.
6
LSR =km s
−1
by averaging over a 6″×6″box centered at
(,
J2000
aJ2000
d
)=[
1
747 20.22,
hms
−28°22′19 7].
b
Present work.
c
T. Hasegawa et al. (2016, in preparation).
d
From
N
dv
CO
18 and X(C
18
O)=
1
10 7
´-(Lis & Goldsmith 1989).
Figure 7. Masses, momenta, and energies of the outflow per unit velocity
width as a function of velocity from the systemic velocity (V
∣
∣D=VLSR–
VSY
S
)
for the blue- and redshifted outflow lobes. The dashed line marks the defined
separation between line core (low velocity)and wing (high-velocity)emission.
9
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

physical parameters of the outflow. The total mass of Sgr B2
(N)outflow is derived as a sum of SiO emission
(
V
LSR =32.2–46.6 km s
−1
for blueshifted velocities and
V
LSR =89.8–109 km s
−1
for redshifted velocities)and SO
2
emission (
V
LSR =48.2–54.6 km s
−1
for blueshifted velocities
and
V
LSR =72.2–85 km s
−1
for redshifted velocities). The
typical outflow velocity,
V
outflo
w
, is derived from the ratio
between momentum and mass. Comparing the spatial structure
of the outflow with other massive outflows (e.g., W51
North;Zapata et al. 2009a), the outflow in Sgr B2(N)has a
relatively symmetric structure.
Ridge & Moore (2001)estimated the parameters of
molecular outflows from 11 high-mass star-forming regions
at a distance of ∼2 kpc. In comparison with their physical
parameters, the Sgr B2(N)outflow has a comparable mass
(∼2000 M), while the flow size and the dynamical timescale
of Sgr B2(N)are an order of magnitude smaller with the other
high-mass star-forming regions. Liu et al. (1998)presented the
SiO outflow from Sgr B2(M)with a mass of 100 Massuming
anX(SiO)of 10
−7
. If they adopt X(SiO)of 10
−8
for Sgr B2
(M), their outflow mass will be similar to the Sgr B2(N)result.
The dynamical timescales of G331.5–01 (Bronfman
et al. 2008; Merello et al. 2013)and W51North (Zapata
et al. 2009a)outflows are comparable to the Sgr B2(N)outflow.
The collimation of the Sgr B2(N)outflow is estimated from
measuring the FWHM of the peak emission of the outflow,
which is ∼0.1 pc from the K2and results in ∼60°. From the
comparisons, the Sgr B2(N)outflow can be characterized as a
young and massivebut moderately collimated outflow.
In comparison with the outflows listed in Zhang et al. (2001)
andBeuther et al. (2002)andthe recompilation of Wu et al.
(2004,2005), the Sgr B2(N)outflows with ∼2000 Msitnear
the massive end of the spectrum. The Sgr B2(N)outflow can be
a single massive outflow expected at the early stages of massive
star formation (e.g., Zapata et al. 2010), or alternatively the
large mass can be a result of overlap and merger of the multiple
outflows as in IRAS16547–4247 (Higuchi et al. 2015).In
order to resolve the outflow completely, observations with
higher spatial resolution are needed.
3.5. Possible Detection of SiO Maser Emission
The maser emission of vibrationally excited SiO in regions
of star formation is rare. It has been detected from only three
regions of massive star formation so far, i.e., Orion-KL (Snyder
& Buhl 1974; Thaddeus et al. 1974), W51 IRS2, and Sgr B2
(M)MD5 (Hasegawa et al. 1986; Ukita et al. 1987; Morita
et al. 1992), despite intensive searches for similar objects in
star-forming regions (Genzel et al. 1980; Barvainis &
Clemens 1984; Jewell et al. 1985; Zapata et al. 2009b).
We checked the ALMA data at the frequencies of the SiO
(J2
1
–=)lines in the v=1, 2, and 3 vibrationally excited
statesand found an emission line for v=2 with a peak flux
density of 2Jy at
V
LSR =72 km s
−1
(Figure 8). The emission
has a line width of FWHM =3.8±0.3 km s
−1
, which is much
narrower than the typical width of thermal emission lines from
this region (FWHM∼10 km s
−1
). Its spatial distribution is
point-like (source size

1 5), and its position, (,
J2000
a
J200
0
d
)=[
17
h
4
7m19. 86,
s−28°22′18 5], is coincident with the
position of K2 within 0 2. The frequency of the detected
emission corresponds to the CH
3
OCHO (v0, 4 3
t2,3,1 1,2,2
=-
)
line at
V
LSR =65 km s
−1
, but this assignment is unlikely
because the emission in other transitions of this molecule hasa
larger line width (FWHM ∼5kms
−1
)and is spatially
extended. The
V
LSR =72 km s
−1
of the SiO (vJ2, 2 1–==
)
line emission is 9 km s
−1
redshifted with respect to the
systemic velocity of
V
LSR =63 km s
−1
, but it is within the
velocity range of the H
2
O masers. This kind of velocity offset
is seen also in W51 IRS2 and Sgr B2(M)MD5 (Hasegawa
et al. 1986; Zapata et al. 2009b). From the total flux of the SiO
(vJ2, 2 1–==
)line emission, the isotropic photon luminos-
ity is estimated as
L
n=2.8 ×10
44
s
−1
. This comfortably falls
within the range of the SiO masers detected in other star-
forming regions (Zapata et al. 2009b). No corresponding signal
was found for the v=1 line exceeding 0.02Jy (3σ). The
spectral range of the v=3 line overlaps with the lines from
CH
2
CH
13
CN, CH
3
CH
3
CO, C
2
H
5
CN (v=1), and CH
3
C(O)
NH
2
, but there seems to be no SiO (v=3, J=2–1)line
stronger than 0.4Jy.
Although SiO masers in various transitions in vibrationally
excited states up to v=4 are detected from many evolved
stars such as Mira variables and red supergiants, the
Table 3
Physical Parameters of the Outflow
Parameter Value
Distance from K2 to SiO peak position
a
:l
1
(pc)
Blue lobe 0.08
Red lobe 0.06
Distance from K2 to SiO outer contour
b
:l
2
(pc)
Blue lobe 0.3
Red lobe 0.2
Outflow mass (
M
)
Blue lobe
c
5.7 ×10
2
Red lobe
d
1.4 ×10
3
Total 2.0 ×10
3
Momentum (
M
km s
−1
)
Blue lobe 7.4 ×10
3
Red lobe 2.2 ×10
4
Total 3.0 ×10
4
Kinetic energy (erg)
Blue lobe 1.1 ×10
48
Red lobe 4.4 ×10
48
Total 5.5 ×10
48
Typical outflow velocity
e
:V
outflow (km s
−1
)
Blue lobe 13
Red lobe 16
Total 15
Dynamical time
f
:
t
dyn (yr)
Blue lobe 5.6 ×10
3
Red lobe 3.7 ×10
3
Average 4.7 ×10
3
Notes. This table shows the physical parameters of the bipolar outflow traced
by SiO and SO
2
; size, mass, momentum, kinetic energy, outflow velocity, and
dynamical time.
a
l
1
is measured as the SiO peak from the K2 position from Figure 5.
b
l
2
is measured as the farthest 5σcontour from the K2 position from Figure 5.
c
Mass, momentum, and kinetic energy of blueshifted components are derived
from integration in the velocity range of V
∣
∣D>7kms
−1
(VLSR =32.2–46.6 km s
−1
for SiO and VLSR =48.2–54.6 km s
−1
for SO
2
).
d
Mass, momentum, and kinetic energy of redshifted components are derived
from integration in the velocity range of V
∣
∣D>7kms
−1
(VLSR =72.2–85 km s
−1
for SO
2
and VLSR =89.8–109 km s
−1
for SiO).
e
V
outflow is derived from the ratio between momentum and mass.
f
t
dyn is calculated from the ratio between l1and V
outflow for the blue- and
redshifted components.
10
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

vJ2, 2 1–==emission is anomalously weakand exhibits
some peculiarity when detected (Clark et al. 1981; Olofsson
et al. 1981,1985; Bujarrabal et al. 1996,2007). Olofsson et al.
(1985)haveproposed that this behavior could arise from a
population transfer due to line overlap between rovibrational
transitions of SiO and H
2
O in the masing regions around
evolved stars. None of the three previously known SiO masers
in star-forming regions havebeen detected in the
vJ2, 2 1–==transition. In the case of the J10–=SiO
maser in W51 IRS2, only the v=2 emission was detected
during 1985–1989, with the v=1 upper limits at 1/5to1/15
of the v=2 intensities (Hasegawa et al. 1986; Fuente
et al. 1989). When Zapata et al. (2009b)observed it in 2003,
they found that the vJ1, 1 0–==emission was detectable at
1.0Jy while the v=2 emission had become weaker at 2.5Jy.
As our understanding of the excitation mechanism for the SiO
masers in star-forming regions is still limited, we cannot rule
out the maser assignment from the fact that only the
vJ2, 2 1–==transition is detected.
The characteristics of the emission line described above, i.e.,
the point-like spatial distribution and extremely narrow line
width, make it quite probable that the line has a maser nature. It
is important to confirm its assignment with the SiO
(vJ2, 2 1–==
)line by, e.g., observing the matching
J10–=lines, measuring polarization, or setting a high enough
lower limit to the line brightness temperature with a higher
spatial resolution. If confirmed, it will be not only the fourth
star-forming region with SiO maser emissionbut also the first
such object with the vJ2, 2 1–==emission. The SiO maser
emission will provide crucial information on the structure,
kinematics, and physical condition of the close vicinity (within
∼100 AU)of the massive protostar that drives the outflow.
3.6. Insight on the Process of Massive Star Formation
Based on the observations of a sample of hot molecular cores
around massive (proto)stars, Cesaroni et al. (2006)and Beltrán
et al. (2011)have proposed that the rotating cores are classified
into two classes, i.e., circumstellar disks and circumcluster
toroids. In their scenario of massive star formation, a larger-
scale infalling envelope provides the mass to the toroid, which
is a transitional structure that feeds the mass toward the
accretion disks with Keplerian rotation around individual
forming stars in the central cluster. Recent observations with
millimeter and submillimeter interferometers provide increas-
ing evidences of disk-like structures of radii 1000–2000 AU
with signatures of Keplerian rotation around B- or O-type
(proto)stars (e.g., Beltrán et al. 2014; Cesaroni et al. 2014;
Hunter et al. 2014; Johnston et al. 2015). Compared with these
cases, the rotating SO
2
core we found around K2 in Sgr B2(N)
may fall in the class of toroids, because it is a much larger ring-
like structure without the signature of the Keplerian rotation.
The structure of Sgr B2(N)presented here shows an
intriguing resemblance to that of W51 North. Both objects
are embedded in very luminous regions of massive star
formationand have the system of bipolar outflow and a large
rotating hot core (e.g., Zapata et al. 2009a,2010). The bipolar
outflow of W51 North is well traced by the SiO (5–4)and CO
(2–1)emission. Its mass is M200 ,~which is an order of
magnitude less than what we find in Sgr B2(N), but it is still a
massive outflow (see, e.g., Zhang et al. 2001; Beuther
et al. 2002; Wu et al. 2004,2005). Near the driving source is
a cluster of luminous H
2
O masers that span a large velocity
rangeand anSiO maser (e.g., Morita et al. 1992; Eisner
et al. 2002). At the center is a peak of millimeter and
submillimeter dust emission, and a hot rotating core surrounds
it (Zapata et al. 2009a,2010). The hot core appears as a rotating
toroid with a central cavity 3000 AU in radius, when observed
in SO
2
(
2
222
2,20 1,21
–)and other molecular lines arising from
energy levels 20–800 K above the ground state. Molecular
emission from even higher energy levels is more spatially
confined and fills in the cavity with the kinematics reproduced
by a model of a Keplerian disk with an infalling motion. Based
on these observations, Zapata et al. (2010)proposed a possible
evolutionary sequence of massive star formation in four phases,
in which W51 North is placed in Phase II, the large and
massive pseudo-disk with layers of physical conditions. Sgr B2
Figure 8. Spectra of molecular lines from the K2 position of Sgr B2(N)around
the frequencies of SiO (v=1, J=2–1)(top, undetected), SiO (v=2, J=2–1)
(middle, tentatively assigned), and SiO (v=3, J=2–1)(bottom, undetected).
The observations were made between 2012 August and October (Belloche
et al. 2014).
11
The Astrophysical Journal, 815:106 (13pp), 2015 December 20 Higuchi et al.

(N)may be in the same phase or a little later, with a continuum
source K2 detected at centimeter wavelengths.
At this point, the structure and kinematics of the molecular
gas in Sgr B2(N)inside the SO
2
toroid are not known. Is it a
single and coherent structure as postulated in the core accretion
model of massive star formation, or, alternatively, a cluster of
forming stars with accretion disks collectively contributing to
the large luminosity and the massive bipolar outflow in the
scenario of competitive accretion (e.g., Tan et al. 2014)? The
modest spatial resolution (∼2″)of the ALMA data we analyzed
here leaves this important question open. Further ALMA
observations with higher spatial resolution would answer the
question and uncover the processes that link the Sgr B2(N)
toroid to the massive star formation inside it.
4. SUMMARY
We have analyzed the archival data of the ALMA
observations of Sgr B2(N)in SO
2
(12 13
4,8 3,11
–)and SiO (2–1)
lines at an angular resolution of ∼2″to investigate the
kinematic structure of the region. Our main findings are
summarized as follows:
1. Sgr B2(N)has a system of a bipolar outflow and a
rotating core, as originally interpreted by Lis et al. (1993).
The SiO line shows the bipolar outflow whose driving
source is located at K2, and the SO
2
line shows both the
outflowing gas and the rotating hot core. We note that
H
2
O masers are associated with the innermost part of the
bipolar outflow.
2. The Sgr B2(N)outflow is characterized as a young
(5×10
3
yr)and massive (∼2000 M)outflow. It is
moderately collimated (∼60°).
3. Fractional abundances of SO
2
and SiO in the outflowing
gas are estimated by a comparison with the C
18
O(1–0)
emission to be X(SO
2
)of
2
.3 0.4
2.
6
-
+×10
−8
and X(SiO)of
1.2 0.1
0.1
-
+×10
−9
for the assumed excitation temperature of
100±50 K. These values are consistent with the
estimates in other high-mass star-forming regions.
4. The mass spectrum of the outflow suggests an approx-
imate power law in the form of
d
Mv()dvµV,
∣
∣Dgwith a
power index 1.
g
~- The shallow slope is consistent
with the youth of the Sgr B2(N)outflow. The outflow
momentum is distributed rather evenly over the velocity
range of the outflow, while a majority of the kinetic
energy is carried at larger V.
∣
∣D
5. We discovered a point source of narrow line emission at
the position of K2 in Sgr B2(N)that can possibly be
assigned to the SiO (vJ2, 2 1–==
)maser emission,
although this assignment needs a confirmation. When
confirmed, this will make Sgr B2(N)the fourth star-
forming region with detected SiO masers.
6. The hot rotating core found in the SO
2
emission has a
ring-like structure with a radius of ∼6000 AU without a
clear sign of Keplerian rotation, and it falls in the class of
toroid classified by Beltrán et al. (2011). Sgr B2(N)
exhibits a striking resemblance to W51 North, although
we do not know the structure inside the toroid of Sgr B2
(N). Compared with W51 North, Sgr B2(N)may be in the
same or a little later phase in the scenario of massive star
formation proposed by Zapata et al. (2010).
Although the overly rich emission lines detected from
Sgr B2(N)require careful selection of molecules and transitions
for proper analyses, ALMA data are proven to be very useful in
understanding the kinematics and physical parameters of this
high-mass star formation region. The Sgr B2 complex is one of
the most important regions as sources for understanding high-
mass star formation in the ALMAera.
We thank the anonymous referee for careful reading and
constructive comments that helped greatly to improve the
manuscript. We also thank the ALMA staff for the observations
during the commissioning stage. This letter makes use of the
following ALMA data: ADS/JAO.ALMA#2011.0.0017.S.
ALMA is a partnership of ESO (representing its member
states), NSF (USA), and NINS (Japan), together with NRC
(Canada), NSC, and ASIAA (Taiwan), in cooperation with the
Republic of Chile. The Joint ALMA Observatory is operated
by ESO, AUI/NRAO, and NAOJ. Data analyses were carried
out on a common-use data analysis computer system at the
Astronomy Data Center, ADC, of the National Astronomical
Observatory of Japan. We acknowledge Takashi Tsukaghi and
Koichiro Sugiyama for their contributions.
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