# Outflow activities in the young high-mass stellar object G23.44-0.18

Jeremy Zhiyuan Ren1, Tie Liu1, Yuefang Wu1, and Lixin Li1,2
1Department of Astronomy, Peking University, 100871, Beijing China, E-mail:
<EMAIL_ADDRESS>
2The Kavli Institute for Astronomy and Astrophysics, Peking University, Yi He
Yuan Lu 5, Hai Dian Qu, Beijing 100871, P. R. China

###### Abstract

We present an observational study towards the young high-mass star forming
region G23.44-0.18 using the Submillimeter Array. Two massive, radio-quiet
dusty cores MM1 and MM2 are observed in 1.3 mm continuum emission and dense
molecular gas tracers including thermal CH3OH, CH3CN, HNCO, SO, and OCS lines.
The 12CO (2–1) line reveals a strong bipolar outflow originated from MM2. The
outflow consists of a low-velocity component with wide-angle quasi-parabolic
shape and a more compact and collimated high-velocity component. The overall
geometry resembles the outflow system observed in the low-mass protostar which
has a jet-driven fast flow and entrained gas shell. The outflow has a
dynamical age of $6\times 10^{3}$ years and a mass ejection rate $\sim
10^{-3}~{}M_{\odot}$ year-1. A prominent shock emission in the outflow is
observed in SO and OCS, and also detected in CH3OH and HNCO. We investigated
the chemistry of MM1, MM2 and the shocked region. The dense core MM2 have
molecular abundances of 3 to 4 times higher than those in MM1. The abundance
excess, we suggest, can be a net effect of the stellar evolution and embedded
shocks in MM2 that calls for further inspection.

###### keywords:

stars: pre-main sequence — ISM: molecules — ISM: abundances — ISM: kinematics
and dynamics — ISM: individual (G23.44-0.18) — stars: formation

††pagerange: Outflow activities in the young high-mass stellar object
G23.44-0.18–LABEL:lastpage††pubyear: 2011

##  1 Introduction

The outflows take place in high-mass star forming regions at very early
evolutionary stages (e.g. Birkmann et al., 2006; Beuther & Sridharan, 2007;
Longmore et al., 2011). They inject large amount of hot gas and kinetic energy
into the molecular cloud, causing intense shock waves that dramatically alter
the chemistry of the surrounding environment (van Dishoeck & Blake, 1998).
However, the morphological and dynamical properties of the outflows on smaller
scales, as well as their chemical effect to the young high-mass stellar cores
are still to be further characterized.

The high-mass star forming region G23.44-0.18 (G23.44 hereafter) was
previously detected as a group of strong CH3OH masers (Walsh et al., 1998)
which indicate the presence of young massive stars. It has a trigonometric-
parallax distance of 5.88 kpc (Brunthaler et al., 2009). In this letter, we
report an observation of the massive dusty cores MM1 and MM2 associated with
the CH3OH masers, as well as a high-velocity, intense bipolar outflow revealed
in CO (2–1). The outflow is causing a shock emission and may be affecting the
chemistry in MM2. Section 3 presents the observational results. A further
discussion on the outflow is given in Section 4.1 and 4.2. The properties of
the dusty cores are discussed in Section 4.3. A summary is given in Section 5.

##  2 Observations and Data Reduction

The observational data is taken from the released SMA data
archive111http://www.cfa.harvard.edu/sma/. The observation toward G23.44 was
carried out in September 2008. The phase tracking center is
RA.(J2000)=18h34m39.25s,
Dec.(J2000)=$-8^{\circ}31^{\prime}36.2^{\prime\prime}$. The pointing accuracy
of the SMA is $\sim 3^{\prime\prime}$. The observation employed 8 antennas in
their compact configuration. The correlator has a total bandwidth of 4 GHz,
centered at 220 GHz (LSB) and 230 GHz (USB). The on-source integration time is
80 min. The SMA primary beam size (field of view) at this waveband (1.3 mm) is
$55^{\prime\prime}$. The correlator has a channel width of 0.812 MHz (1.1 km
s-1). The system temperature is between 120 to 150 K. QSO 3c454.3 and Neptune
were taken as bandpass and flux calibrators, respectively. QSO 1733-130 and
1911-201 were observed in every 30 and 50 minutes respectively during the
observation sequence to track the antenna gains. The interferometer array has
a shortest baseline of 17 m, corresponding to an u-v coverage for structures
smaller than $20^{\prime\prime}$.

The calibration and imaging were performed in Miriad1. The absolute flux level
has an uncertainty of $\sim 10\%$. The continuum map was extracted from the
line-free channels in LSB. As the visibility data are INVERTed to the image
domain, the synthesized beam size is $4.1^{\prime\prime}\times
3.7^{\prime\prime}$.

Images of the G23.44 region at other wavelengths were also extracted from
several public archives. The Spitzer/IRAC images and the point-source
catalogue are taken the database of the GLIMPSE sky survey in the NASA/IPAC
Infrared Science Archive222http://irsa.ipac.caltech.edu/. We looked for the
near-infrared counterparts in the 2MASS point-source catalogue2. The continuum
image at JCMT 450 $\micron$ was also obtained to measure the flux densities.
They are taken from the Canadian Astronomy Data Center (CADC) repository of
the SCUBA Legacy Fundamental Object Catalogue333http://www4.cadc-ccda.hia-
iha.nrc-cnrc.gc.ca.

Figure 1: The 1.3 mm continuum emission observed with the SMA overlayed on the
IRAC composite image (blue=3.6, green=4.5, and red=8.0 $\micron$) image. The
image center is RA.(J2000)=18h34m39.25s and
Dec.(J2000)=$-8^{\circ}31^{\prime}36.2^{\prime\prime}$. The contour levels are
-5, 5, 15, 25, 35, 45, and 55 $\sigma$ (0.013 Jy beam-1). The synthesized beam
($4.1^{\prime\prime}\times 3.7^{\prime\prime}$) with $PA=7^{\circ}$ from the
north to NW. The squares denote the two strongest CH3OH masers. The crosses
are the centers of the IRAC sources. Table 1: Physical parameters of the
cores.

parameter | MM1 | MM2 | Unit
---|---|---|---
$m_{J}~{}^{a}$ | $12.315\pm 0.049$ | $>15.8$ | mag
$m_{H}~{}^{a}$ | $11.702\pm 0.028$ | $>15.1$ | …
$m_{K_{s}}~{}^{a}$ | $11.435\pm 0.023$ | $>14.3$ | …
$F(3.6\micron)^{b}$ | $1.42\pm 0.28$ | $0.63\pm 0.08$ | mJy
$F(4.5\micron)^{b}$ | $25.65\pm 2.23$ | $14.48\pm 1.96$ | …
$F(5.8\micron)^{b}$ | $47.84\pm 2.22$ | $53.16\pm 2.11$ | …
$F(8.0\micron)^{b}$ | $32.92\pm 2.04$ | $40.64\pm 2.86$ | …
$F(450\micron)^{c}$ | $27.3\pm 1.5$ | $38.6\pm 1.5$ | Jy
F(1.3 mm) | $0.65\pm 0.03$ | $1.42\pm 0.03$ | …
$T_{\rm rot}$(CH3CN) | $65\pm 10$ | $110\pm 10$ | K
$T_{\rm rot}$(CH3OH) | $66\pm 5$ | $60\pm 6$ | …
Mass | $120\pm 30$ | $140\pm 10$ | $M_{\odot}$
$N({\rm H_{2}})$ | $1.9\pm 0.2$ | $2.5\pm 0.2$ |  $10^{23}$ cm-2

Note. The flux densities at IR and sub-millimeter wavebands are taken from
a2MASS, bSpitzer/IRAC, cJCMT/SCUBA. The remaining data all come from the SMA
observation.

##  3 Results

Figure 1 shows the 1.3 mm continuum emission (white contours) overlayed on the
RGB image of the three IRAC bands. It reveals two dust cores aligned from the
north to south, denoted as MM1 (northern core) and MM2 (southern core).
Deconvoluted with the beam size, MM1 has an mean diameter of
$7.3^{\prime\prime}$ (0.2 pc) while MM2 is more elliptical, with an extent of
$9.7^{\prime\prime}\times 6.7^{\prime\prime}$ ($0.27\times 0.19$ pc) and
position angle $PA=45^{\circ}$ northeast. The strongest 6.7 GHz CH3OH masers
(Walsh et al., 1998) coincide with the centers of MM1 and MM2. We referred to
the MAGPIS 6 cm survey444http://third.ucllnl.org/gps/index.html for the
potential radio continuum emission from the ionized gas. It turns out that
neither MM1 nor MM2 has detectable emission above the sensitivity level of 2.9
mJy, indicating the both cores being prior to forming an Ultra-Compact Hii(UC
Hii) region. The fluxes (or magnitudes) of MM1 and MM2 at near- to far-
Infrared are drawn from the 2MASS and IRAC point-source catalogues. The 450
$\micron$ fluxes are measured directly from the JCMT images. The obtained
values are shown in Table 1.

Figure 2: (a) The velocity-integrated images of the molecular lines (contours)
overlayed on the 1.3 mm continuum (gray). The integration is from 95 to 110 km
s-1. For C18O (2–1), the contours are -4, 4, 8, 12… 40 $\sigma$ ($1.35$ Jy
beam-1 km s-1). For CH3CN $(12_{2}-11_{2})$, the contours are -4, 4, 8… 24
$\sigma$ ($0.6$ Jy beam-1 km s-1). For HNCO ($10_{0,10}-9_{0,9}$), the
contours are -4, 4, 6, 8… 20 $\sigma$ ($0.34$ Jy beam-1 km s-1). For CH3OH
$(8_{0}-7_{1})$, the contours are -4, 4, 8… 48 $\sigma$ ($0.5$ Jy beam-1 km
s-1). For SO ($6_{5}-5_{4}$), the contours are -4, 4, 8, 12… 32 $\sigma$
($0.6$ Jy beam-1 km s-1). For OCS (19–18), the contours are -4, 4, 5… 11
$\sigma$ ($0.3$ Jy beam-1 km s-1). The cross denotes the position of the
shocked clump at
$(\Delta\alpha=-11^{\prime\prime},\Delta\delta=10^{\prime\prime})$. (b) The
components of the 12CO (2–1) outflow. The low velocity component (LVC) is
integrated in (83,93) and (113,123) km s-1 for its blue- and redshifted lobes,
respectively, while the high velocity component (HVC) is integrated in (65,75)
and (140,170) km s-1 for the two lobes. For each lobe, the contours are -10,
10, 20, …, 90% of the peak value. The dashed line denotes the axis of the
outflow as indicated by the HVC. (c) The beam-averaged SO and OCS spectra at
the center of MM1 and MM2, and the shocked clump. The vertical lines denote
the systemic velocity of $v_{\rm lsr}=101$ km s-1.

The $J=2-1$ transition of 12CO, 13CO and C18O and a number of molecular
species tracing dense gas are detected in our sidebands, including HNCO
($10_{0,10}-9_{0,9}$), SO ($6_{5}-5_{4}$), OCS (19–18), CH3CN $12_{K}-11_{K}$
($K=0$ to 6), and four thermally excited CH3OH lines. In the online material,
their spectra and physical parameters are provided in Figure S1 and Table S1.
The 2MASS $K_{s}$ and JCMT 450 $\micron$ images are presented in Figure S2 and
S3.

Figure 2a shows the integrated images of the C18O, CH3CN, HNCO, CH3OH
$(8_{0}-7_{1})$, SO, and OCS lines. The C18O has a broad distribution, with
the emission peaks reasonably coincident with the dust cores. The other
species show more compact emissions closely associated with the dust cores.
The C18O emission has comparable intensities for MM1 and MM2, while the
emissions of other dense-gas tracers are much brighter in MM2. Since the dust
continuum and C18O emission all closely follow the H2 distributions, the
intensity contrast suggests that MM2 have higher molecular abundances than
MM1, as discussed in Section 4.3.

A strong bipolar outflow is observed in 12CO (2–1). Figure 2b presents the CO
(2–1) image integrated in the four velocity intervals. The outflow can be
divided into a pair of low-velocity component (LVC) and high-velocity
component (HVC) based on their distinct morphologies. The HVC roughly has
$|v|\geq 20$ km s-1 from the systemic velocity (101 km s-1). The bulk of the
outflow emission is in the southeast-northwest direction, with a position
angle of $PA=40^{\circ}$ northwest. Figure 2c shows the beam-averaged SO and
OCS spectra at MM1 and MM2 center, and the shocked region
$(\Delta\alpha=-11^{\prime\prime},\Delta\delta=10^{\prime\prime})$. The 13CO
emission shows a similar feature with 12CO despite being weak in the HVC. We
present the channel maps of 12CO and 13CO in Figure S4 and S5, respectively.

##  4 discussion

###  4.1 The physical properties of the outflow

The presence of intense bipolar outflow in G23.44 suggests that the central
stars are likely being formed via a disk-mediated accretion (Zhang et al.,
2007). The HVC is well collimated along the outflow axis, while the LVC,
especially its red wing, has a more opened parabolic shape with a decreasing
brightness from the edge to the center. The overall geometry of the outflow is
reminiscent of the outflow system in the low-mass protostar HH211 (Gueth &
Guilloteau, 1999). HH211 consists of a high-velocity jet-driven flow and a
entrained low-velocity gas shell. The G23.44 outflow may represent a scaled-up
version of this system, despite its collimation and shell structure being much
less perfect than HH211. However, the size of the outflow in HH211 is only
$\sim 10^{4}$ AU, while the G23.44 outflow ($10^{5}$ AU) has a 10-time larger
spatial extent. On such a large scale, the outflow structure might begin to
deteriorate due to increased turbulence in its environment.

Table 2: Abundance of the detected molecules.

species | $N_{\rm x}$ | $f_{\rm x}$ | $f_{x,2/1}^{a}$
---|---|---|---
(X) |  (cm-2) |  |
MM1 $(-1^{\prime\prime},10^{\prime\prime})$
C18O | $3.16\pm 0.05(16)$ | $1.7\pm 0.1(-7)$ | $-$
CH3OH | $2.5\pm 1.2(15)$ | $1.3\pm 0.6(-8)$ | $-$
CH3CN | $5.2\pm 0.6(14)$ | $2.7\pm 0.5(-9)$ | $-$
HNCO | $3.8\pm 0.2(14)$ | $2.0\pm 0.4(-9)$ | $-$
SO | $2.5\pm 0.2(14)$ | $1.3\pm 0.3(-9)$ | $-$
OCS | $3.7\pm 0.2(14)$ | $1.9\pm 0.3(-9)$ | $-$
MM2 $(1^{\prime\prime},-3^{\prime\prime})$
C18O | $7.30\pm 0.07(16)$ | $2.9\pm 0.1(-7)$ | $1.7\pm 0.2$
CH3OH | $9.8\pm 0.8(15)$ | $3.9\pm 0.9(-8)$ | $3.0\pm 1.5$
CH3CN | $1.6\pm 0.2(15)$ | $6.3\pm 0.7(-9)$ | $2.3\pm 0.8$
HNCO | $2.20\pm 0.05(15)$ | $8.8\pm 1.5(-9)$ | $4.4\pm 1.4$
SO | $1.45\pm 0.06(15)$ | $5.8\pm 0.9(-9)$ | $4.5\pm 1.4$
OCS | $1.20\pm 0.05(15)$ | $4.8\pm 0.8(-9)$ | $2.5\pm 0.7$
outflow-shock $(-11^{\prime\prime},10^{\prime\prime})$
C18O | $2.4\pm 0.07(16)$ | $=2.0(-7)^{b}$ | $-$
CH3OH | $2.8\pm 0.5(15)$ | $2.3\pm 0.5(-8)$ | $-$
HNCO | $4.1\pm 0.2(14)$ | $8.3\pm 0.4(-10)$ | $-$
SO | $3.5\pm 0.3(14)$ | $3.1\pm 0.3(-9)$ | $-$
OCS | $8.0\pm 0.3(14)$ | $6.6\pm 0.3(-9)$ | $-$

$a.$ The abundance ratio of MM2 to MM1.

$b.$ Assumed as the typical value in ISM. Abundances of the other molecules
are then calculated by comparing to $N({\rm C^{18}O})$.

The CO outflow, especially the LVC also exhibits a weak emission feature in
opposite direction to the major one, i.e. a blueshifted emission in the NW and
a redshifted one in the SE. This can be a second flow by chance aligned in the
same position angle with the major one. Alternatively, once the outflow
direction is close to the plane of the sky and its inclination angle becomes
smaller than the opening angle $(\sim 30^{\circ})$, its front and back sides
will exhibit opposite Doppler shift. This scenario is well demonstrated by Yen
et al. (2010, Figure 10 therein). We speculate the second case to be more
plausible.

Assuming a local thermal equilibrium (LTE) and low optical depth for the 12CO
line-wing emission, the CO column density and the outflow mass can be inferred
following the approach of Garden et al. (1991). The excitation temperature of
the outflow is speculated to be slightly lower than in the dense core MM2
($\sim 70$ K, Section 4.3). A value of $T_{\rm ex}=50$ K is adopted here. The
outflow mass is estimated in each 1 km s-1 interval, then the total mass,
momentum, and kinetic energy are calculated from $M=\Sigma_{v}m(v)$,
$P=\Sigma_{v}m(v)v$, and $E_{k}=\Sigma_{v}m(v)v^{2}/2$. In calculation, we
assume an inclination angle of $\theta=20^{\circ}$ to correct the velocity
projection. The outflow turns out to have $M\simeq 5.5~{}M_{\odot}$, $P\simeq
360~{}M_{\odot}$ km s-1, $E_{k}\simeq 9\times 10^{46}$ erg.

Each of the blue and red wings in the HVC is resolved into two major clumps
(Figure 2b, lower panel) denoted B1, B2 and R1, R2. These clumps in the
outflow may indicate an episodic mass ejection (e.g. Qiu et al., 2009). The
average central distance is $24^{\prime\prime}$ (0.68 pc) for the pair of
B1-R1, and $5.0^{\prime\prime}$ (0.14 pc) for B2-R2. The two pairs both show a
reasonable geometrical symmetry. However, the clumps are not exactly aligned
across the MM2 center, instead with an offset to its west. This may indicate
that the driven agent, or the stellar disk system is also displaced from the
MM2 center. The dynamical age of the outflow $t_{\rm dyn}$ is estimated by
dividing its spatial extent with the typical velocity ($v=40/\sin\theta$ km
s-1). That yields $t_{\rm dyn}=6.1\times 10^{3}$ years for the B1-R1 pair and
$t_{\rm dyn}=1.3\times 10^{3}$ years for B2-R2. Adopting the first value, we
calculate the mass ejection rate to be $\dot{M}=M/t_{\rm dyn}\simeq 1.0\times
10^{-3}M_{\odot}$ yr-1. One can see that the outflow is young and performing
an intense mass ejection.

###  4.2 Chemical enhancement in the outflow and shock

In both low- and high-mass young stellar objects the outflows have led to a
rich chemistry in CH3OH, HNCO, and sulphides (Blake et al., 1987; van der Tak
et al., 2003; Jørgensen et al., 2004; Rodríguez-Fernández et al., 2010, etc.).
In G23.44, the SO and OCS emission shows a noticeable emission clump at
offset=$(-11^{\prime\prime},10^{\prime\prime})$ (Figure 2a). The clump is
almost perfectly aligned with the HVC axis of the CO (2–1), and reasonably
coincides with the peak R2, despite being closer to MM2 center. This strongly
suggests that the clump traces a shock within the jet/outflow. The clump is
also weakly detected in HNCO and CH3OH (spectra shown in Figure S1). Their
column densities are estimated using Equation (1) and (2) in Section 4.3. In
calculation a temperature of $T_{\rm ex}=50$ K [same with that of CO (2–1)] is
adopted, and the C18O abundance is taken to be a constant of $2\times 10^{-7}$
as the typical ISM abundance which is close to $f_{\rm C^{18}O}$ measured in
MM1 and MM2. The abundances of other molecules are then inferred from their
intensity ratio relative to C18O (2–1). The abundances of the shocked clump
are listed in Table 2.

In the low-mass protostars, the outflows often lead to a rich chemistry. In a
few cases, the column density of HNCO and S-bearing molecules can been
enhanced for orders of magnitude (e.g. L1157 and L1448, Rodríguez-Fernández et
al., 2010; Taffala et al., 2010). In comparison, the shocked region at
$(-11^{\prime\prime},10^{\prime\prime})$ shows moderate abundances which are
comparable or even slightly lower than in MM2. The enhancement level can be
limited by the temperature and density in the shocked region, as well as the
efficiency of the previous grain-surface chemistry. In addition, since the
outflow has a short $t_{\rm dyn}$, the shock chemistry may still be undergoing
and have not yet attained its maximum level.

The blue lobe of CO $(2-1)$ outflow are devoid of the SO and OCS emissions. As
a possible explanation, the S-bearing species (e.g. H2S and CS as their
progenitors) would be rather tightly bound to the dust grains (Blake et al.,
1987), and and a threshold shock velocity may be required for their
desorption. Higher resolution and primary shock products like SiO and H2S may
better reveal the shock chemistry.

Figure 3: (a) The solid line is the CH3CN ($12_{K}-11_{K}$) towards the center
of MM1 (top panel) and MM2 (bottom panel). The dashed line is the spectrum of
the best fit from the LTE gas model. (b) The rotation diagram of the CH3OH
lines. The filled circles are for the observed transitions. The vertical bars
indicate the $3\sigma$ errors. The linear least-squares fit is shown as the
solid line.

###  4.3 The dense core properties

The dust cores MM1 and MM2 coincide with dark patches on the extended 8
$\micron$ emission (Figure 1). The two cores are slightly connected with each
other, suggesting that they might be the fragmentation products from the same
natal cloud.

The CH3CN ($12_{K}-11_{K}$) lines provide a temperature diagnosis for the
cores (Figure 3a). We take an LTE radiative transfer model with a uniform
rotation temperature and gas density (Chen et al., 2006; Zhang et al., 2007)
to reproduce the observed CH3CN spectra at the dust-core center. We found the
best fit at $T_{\rm rot}=65$ K for MM1 and $T_{\rm rot}=110$ K for MM2. The
modeled spectra of $K=0$ and 1 towards MM2 is much lower than the observed
line profile while the $K>1$ lines are all well fitted. The excess emission in
$K=0$ and 1 lines may largely arise from the cold outer-envelope gas.

The rotation temperature is also evaluated from the thermally excited CH3OH
lines at the core center using the rotation diagram method (Goldsmith &
Langer, 1999; Liu et al., 2001; Bisschop et al., 2007). The column density of
the upper level $N_{\rm u}$ is related to the line intensity in the form

$\qquad\qquad\frac{N_{\rm u}}{g_{\rm u}}=\frac{3k}{8\pi^{3}\nu S\mu^{2}}\int
T_{\rm b}\,{\rm d}v$ (1)

where $g_{\rm u}$ is the upper-level degeneracy. $S$ is the line strength,
$\mu$ is the dipole moment. $N_{\rm u}$ depends on the total column density
$N_{\rm T}$ in the form

$\qquad\qquad\ln(\frac{N_{\rm u}}{g_{\rm u}})=\ln(\frac{N_{\rm T}}{Q_{\rm
rot}})-\frac{E_{\rm u}}{T_{\rm rot}}$ (2)

where the partition function $Q_{\rm rot}$ can be adopted from Blake et al.
(1987). $T_{\rm rot}$ is then estimated from a least-square fit (Figure 3b).
The derived values are listed in Table 1.

The $T_{\rm rot}$ derived from CH3OH and CH3CN are remarkably similar for MM1.
For MM2, however, the $T_{\rm rot}$ of CH3OH is lower than that of CH3CN. As a
probable reason, the CH3OH lines generally have lower $E_{\rm u}$ than CH3CN
($12_{K}-11_{K}$), then the cooler outer-part gas would contribute more
significantly to the CH3OH lines than to the CH3CN (in particular $K=2$ to 6).
In addition the CH3OH can be easily released from the dust grains by the
shocks while the CH3CN production requires subsequent gas-phase reactions and
stellar heating (van Dishoeck & Blake, 1998). CH3CN may therefore be
distributed closer to the MM2 center (as seen in Figure 2a), and showing a
higher $T_{\rm rot}$.

For the other species, an independent temperature estimation is currently
unavailable. We adopt $T_{\rm ex}=60$ K for MM1 and $T_{\rm ex}=70$ K for MM2
assuming an LTE excitation, the column density of the molecules are calculated
from their spectral lines towards the dust-core center, using equation (1) and
(2).

Assuming optically thin 1.3 mm dust continuum emission, the total dust-and-gas
mass can be estimated using $S_{\nu}=M_{\rm core}\kappa_{\nu}B_{\nu}(T_{\rm
d})/gD^{2}$. In the formula $S_{\nu}$ is the total flux density at 1.3 mm.
$g=100$ is the dust-gas mass ratio. $B_{\nu}(T_{\rm
d})=(2h\nu^{3}/c^{2})/[\exp(T_{0}/T_{\rm d})-1]$ is the Planck function
($T_{0}=h\nu/k$). The dust opacity is set to be $\kappa_{\nu}=\kappa_{\rm
230GHz}=0.9$ cm2 g-1 of MRN dust grains with thin ice mantles at $n_{\rm
H_{2}}=10^{6}$ cm-3 (Ossenkopf & Henning, 1994). In the expression of
$S_{\nu}$, using the beam-averaged flux density towards the core center, and
dividing the obtained mass with $m_{\rm H_{2}}$, one can also calculate the
peak column density N(H2) at MM1 and MM2. The derived values are presented in
Table 1. MM2 has a comparable but slightly higher N(H2) than MM1.

The molecular abundances are estimated using $f_{\rm x}=N(X)/N({\rm H_{2}})$.
The $f_{\rm x}$ and its ratio between MM2 and MM1 are listed in Table 2. For
the high-density tracers, the abundances in MM2 are generally 3 to 4 times
higher than those in MM1. The SO shows the highest abundance ratio of 4.5. We
note that the temperature adopted for MM2 (70 K) is conservative. A higher
$T_{\rm ex}$ would yield a lower N(H2) and higher $N(X)$ for most species,
hence leading to even higher abundances. For instance, adopting the value of
$T_{\rm rot}({\rm CH_{3}CN})=110$ K as the $T_{\rm ex}$ for MM2, we will have
$N({\rm HNCO})=6.7\times 10^{15}$ cm2 that is 3 times larger than the current
value.

The molecular distribution appears to be influenced by the outflow. Besides
the coherent elongation of SO and OCS with the HVC of 12CO emission (Figure
2b), the CH3OH and HNCO emissions are also largely distributed to the west of
MM2. In addition, all the molecular lines have broader line widths in MM2 than
in MM1, e.g. The SO line has $\Delta v_{\rm FWHM}=11.3$ in MM2 and 3.4 km s-1
in MM1. The line widths of MM2 can be broadened by a kinetic energy input from
the outflow. Considering its influence to the molecular line profiles, the
outflow may also contribute to the higher abundances in MM2. Specifically,
there would be underlying shocks which are induced as the outflow is
interacting with the envelope and/or infalling gas (Arce & Sargent, 2004; Wu
et al., 2009).

At the near- to mid-IR and sub-millimeter wavebands, the two cores have
similar flux densities (Table 1). At IRAC bands MM2 is brighter than MM1 at
5.8 and 8 $\micron$ but weaker at 3.6 and 4.5 $\micron$, while at 2MASS $J$,
$H$, $K_{s}$ bands, only MM1 is significantly detected (Figure S2). Having a
spectral energy distribution (SED) obviously skewed to the shorter wavelengths
($\lambda<4.5~{}\micron$), MM1 may have a more evolved stage for its stellar
disk system (Robitaille et al., 2006). However, since the outflow in MM2 is
close to the plane of the sky, it is also possible that a disk/toroid
structure lies edge-on in MM2 thus would largely absorb the short-wave
emissions. A higher resolution will be helpful in revealing the morphology of
the stellar system and underlying shocks in MM2.

##  5 summary

We present a multi-wavelength study towards the young high-mass star forming
region G23.44-0.18. Two dusty cores, MM1 and MM2 are observed in 1.3 mm
continuum emission and the molecular lines of CH3OH, CH3CN, HNCO, SO, and OCS.
The both cores have a pre-UC Hii evolutionary stage. A strong bipolar outflow
is arising from MM2 as revealed by CO (2–1). The outflow consists of a
collimated, bi-polar high-velocity component and a more extended, parabolic
flow at lower velocities. A clump of shocked gas is observed in SO
$(6_{5}-5_{4})$ and OCS $(19-18)$ lines and also detected in HNCO and CH3OH.
The outflow shows a high momentum and mass-loss rate, and a short timescale of
$t_{\rm dyn}\sim 6\times 10^{3}$ years.

In MM2 the high-density tracers CH3OH, CH3CN, HNCO, SO, and OCS show
abundances of 3 to 4 times higher than those in MM1. To explain this
difference, we suggest a combined effect of the stellar heating and underlying
shock interactions in MM2. A further study with higher excited lines and
improved resolution can be performed to probe the chemistry, possible disk-jet
system and mass accretion close to the MM2 center.

## Acknowledgment

We are grateful to the SMA observers and the SMA data archive. We would thank
the anonymous referee for the useful comments that helped to improve the
presentation and interpretation. This work is supported by grants of
No.10733033, 10873019, 10973003, 2009CB24901, and the Doctoral Candidate
Innovation Research Support Program (kjdb201001-1) from Science & Technology
Review.

## References

  * Arce & Sargent (2004) Arce, H. & Sargent, A. 2004, ApJ, 612, 342
  * Beuther & Sridharan (2007) Beuther, H., & Sridharan, T. K. 2007, ApJ, 668, 348
  * Birkmann et al. (2006) Birkmann, S. M., Krause, O., Lemke, D. 2006, ApJ, 637, 380
  * Bisschop et al. (2007) Bisschop, S. E., Jorgensen, J. K., van Dishoeck, E. F., & de Wachter, E. B. 2007, A&A, 465, 913
  * Blake et al. (1987) Blake, G., Sutton, E., Masson, C., & Phillips, T. 1987, ApJ, 315, 621
  * Brunthaler et al. (2009) Brunthaler, A., Reid, M. J., Menten, K. M., Zheng, X. W., Moscadelli, L., & Xu, Y. 2009, ApJ, 693, 424
  * Chen et al. (2006) Chen, H., Welch, W. J., Wilner, D. J., Sutton, E. C. 2006, ApJ, 639, 975
  * Garden et al. (1991) Garden, P. R., Hayashi, M., Gatley, I. et al., 1991, ApJ, 374, 540
  * Goldsmith & Langer (1999) Goldsmith P. & Langer, W. 1999, ApJ, 517, 209
  * Gueth & Guilloteau (1999) Gueth, F. & Guilloteau, S. 1999, A&A, 343, 571
  * Jørgensen et al. (2004) Jørgensen, J. K., Hogerheijde, M. R., Blake, G. A., van Dishoeck, E., Mundy, L., & Schöier, F., A&A, 415, 1021
  * Liu et al. (2001) Liu, S-Y., Mehringer, D., & Snyder, L. 2001, ApJ, 552, 654
  * Longmore et al. (2011) Longmore S. N., Pillai, T., Keto, E., Zhang, Q., and Qiu, K. 2011, ApJ, 726, 97
  * Ossenkopf & Henning (1994) Ossenkopf, V., & Henning, T. 1994, A&A, 291, 943
  * Qiu et al. (2009) Qiu, K., Zhang, Q., Wu, J., & Chen, H. 2009, ApJ, 696, 66
  * Rodríguez-Fernández et al. (2010) Rodríguez-Fernández, N., Tafalla, M., Gueth, F., & Bachiller, R. 2010, A&A, 516, 98
  * Robitaille et al. (2006) Robitaille, T. P., Whitney, B. A., Indebetouw, R., Wood, K., & Denzmore, P. 2006, ApJS, 167, 256
  * van der Tak et al. (2003) van der Tak, Boonman, A., Braakman, R., van Dishoeck, E. 2003, A&A, 412, 133
  * Taffala et al. (2010) Tafalla, M., Santiago-García, J., Hacar, A., & Bachiller, R. 2010, A&A, 522, 91
  * van Dishoeck & Blake (1998) van Dishoek, E. & Blake, G. 1998, ARA&A, 36, 317
  * Walsh et al. (1998) Walsh, A. J., Burton, M. G., Hyland, A. R., Robinson, G. 1998, MNRAS, 301, 640
  * Wu et al. (2009) Wu, Y., Qin, S. L., Guan, X., Xue, R., Ren, Z., Liu, T., Huang, M., Chen, S. 2009, ApJ, 697, L116
  * Yen et al. (2010) Yen, H., Takakuwa, S., Ohashi, N. 2010, ApJ, 710, 1786
  * Zhang et al. (2007) Zhang, Q., Hunter, T. R., Beuther, H., Sridharan, T. K., Liu, S.-Y., Su, Y.-N., Chen, H.-R., & Chen, Y. 2007, ApJ, 658, 1152

Table S1: Parameters of the molecular lines. Frequency | Species |  | $V_{\rm lsr}$ | $T_{\rm b}^{\rm peak}$ | $\Delta V_{\rm FWHM}$
---|---|---|---|---|---
(GHz) | (X) | Transition |  (km s-1) | (K) |  (km s-1)
MM1 $(-1^{\prime\prime},10^{\prime\prime})$
219.5603… |  C18O | $2-1$ | 102.2 | $11\pm 0.4$ | $3.2\pm 0.5$
219.7982… | HNCO | $10_{0,10}-9_{0,9}$ | 101.4 | $1.8\pm 0.1$ | $3.6\pm 0.4$
219.9494… | SO | $6_{5}-5_{4}$ | 101.3 | $4.2\pm 0.1$ | $3.4\pm 0.4$
220.0784… |  CH3OH | $8_{0,8}-7_{1,6}$ | 101.5 | $1.0\pm 0.3$ | $4.2\pm 0.6$
220.7302… |  CH3CN | $12_{2}-11_{2}$ | 101.6 | $1.3\pm 0.2$ | $5.5\pm 0.8$
229.7588… |  CH3OH | $8_{-1}-7_{0}$ | 101.8 | $3.9\pm 0.1$ | $3.5\pm 0.5$
230.0270… |  CH3OH | $3_{-2,2}-4_{-1,4}$ | 102.1 | $1.3\pm 0.3$ | $2.9\pm 0.6$
231.0609… | OCS | $19-18$ | 102.1 | $1.9\pm 0.2$ | $4.5\pm 0.4$
231.2810… |  CH3OH | $10_{2,9}-9_{3,6}$ | 101.3 | $0.6\pm 0.2$ | $2.3\pm 0.4$
MM2 $(1^{\prime\prime},-3^{\prime\prime})$
219.5603… |  C18O | $2-1$ | 99.8 | $6.4\pm 0.5$ | $10.5\pm 1.5$
219.7982… | HNCO | $10_{0,10}-9_{0,9}$ | 100.6 | $3.5\pm 0.1$ | $8.8\pm 0.6$
219.9494… | SO | $6_{5}-5_{4}$ | 101.5 | $6.0\pm 0.1$ | $11.3\pm 0.7$
220.0784… |  CH3OH | $8_{0,8}-7_{1,6}$ | 100.5 | $3.5\pm 0.3$ | $6.7\pm 0.8$
220.7302… |  CH3CN | $12_{3}-11_{3}$ | 101.7 | $2.8\pm 0.1$ | $6.6\pm 0.7$
229.7588… |  CH3OH | $8_{-1}-7_{0}$ | 102.1 | $15.0\pm 0.1$ | $8.5\pm 0.8$
230.0270… |  CH3OH | $3_{-2,2}-4_{-1,4}$ | 102.1 | $2.0\pm 0.2$ | $4.3\pm 0.5$
231.0609… | OCS | $19-18$ | 102.1 | $3.2\pm 0.1$ | $6.7\pm 0.4$
231.2810… |  CH3OH | $10_{2,9}-9_{3,6}$ | 102.1 | $1.4\pm 0.3$ | $4.3\pm 0.3$
Shock $(-11^{\prime\prime},10^{\prime\prime})$
219.5603… |  C18O | $2-1$ | 102.1 | $8.0\pm 0.3$ | $2.7\pm 0.4$
219.7982… | HNCO | $10_{0,10}-9_{0,9}$ | 96.3 | $0.8\pm 0.2$ | $1.3\pm 0.4$
219.9494… | SO | $6_{5}-5_{4}$ | 102.3 | $4.9\pm 0.3$ | $5.6\pm 0.5$
220.0784… |  CH3OH | $8_{0,8}-7_{1,6}$ | 102.1 | $0.9\pm 0.2$ | $3.7\pm 0.6$
220.7302… |  CH3CN | $12_{2}-11_{2}$ | $-$ | $-$ | $-$
229.7588… |  CH3OH | $8_{-1}-7_{0}$ | 103.2 | $8.0\pm 0.2$ | $4.2\pm 0.4$
230.0270… |  CH3OH | $3_{-2,2}-4_{-1,4}$ | $-$ | $-$ | $-$
231.0609… | OCS | $19-18$ | 102.1 | $2.0\pm 0.2$ | $3.3\pm 0.5$
231.2810… |  CH3OH | $10_{2,9}-9_{3,6}$ | $-$ | $-$ | $-$

Figure S1: The molecular lines detected in the observational sidebands at the
three positions.
Figure S2: The near to mid-IR RGB-coded image of G23.44. The blue is the 2MASS
$K_{s}$, the green and red colors are IRAC 3.6 and 4.5 micron, respectively.
The image center is RA.(J2000)=18h34m39.25s and
Dec.(J2000)=$-8^{\circ}31^{\prime}36.2^{\prime\prime}$. The contour levels are
-5, 5, 15, 25, 35, 45, and 55 $\sigma$ (0.013 Jy beam-1). The synthesized beam
($4.1^{\prime\prime}\times 3.7^{\prime\prime}$) with $PA=7^{\circ}$ to the
west. The squares are the two strongest CH3OH masers. The crosses are the
centers of the IRAC sources.
Figure S3: The SCUBA 450 $\micron$ image of G23.44. Contours are 10 to 90 per
cent of the peak brightness (34.5 Jy beam-1). The beam size is
$7.5^{\prime\prime}$. The cross denotes the center position
RA.(J2000)=18h34m39.25s and
Dec.(J2000)=$-8^{\circ}31^{\prime}36.2^{\prime\prime}$.

Figure S4: The 12CO (2–1) channel images. The contours are -10, 20 to 90
percent of the peak intensity, which is 45 Jy beam-1 for the velocity range of
[63,90] km s-1 and 18.9 Jy beam-1 for v=[110,152] km s-1. Gray is the SMA 1.3
mm continuum emission.
Figure S5: The 13CO (2–1) channel images. The contours are -10, 20 to 90
percent of the peak intensity which is 12.2 Jy beam-1 for the velocity range
of [84,93] km s-1 and 8.8 Jy beam-1 for v=[108,117] km s-1.