Astronomy & Astrophysics manuscript no. beuther˙g29 c© ESO 2018
November 4, 2018
Dust and gas emission in the prototypical hot core G29.96−0.02
at sub-arcsecond resolution
H. Beuther1, Q. Zhang2, E.A. Bergin3, T.K. Sridharan2, T.R. Hunter4, S. Leurini5
1 Max-Planck-Institute for Astronomy, Königstuhl 17, 69117 Heidelberg, Germany
e-mail: beuther@mpia.de
2 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA
e-mail: name@cfa.harvard.edu
3 University of Michigan, Dept. of Astronomy, Ann Arbor, MI 48109-1090
e-mail: ebergin@umich.edu
4 NRAO, 520 Edgemont Rd, Charlottesville, VA 22903
e-mail: thunter@nrao.edu
5 European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany
e-mail: sleurini@eso.org
ABSTRACT
Context. Hot molecular cores are an early manifestation of massive star formation where the molecular gas is heated to temperatures > 100 K
undergoing a complex chemistry.
Aims. One wants to better understand the physical and chemical processes in this early evolutionary stage.
Methods. We selected the prototypical hot molecular core G29.96−0.02 being located at the head of the associated ultracompact Hii region.
The 862 µm submm continuum and spectral line data were obtained with the Submillimeter Array (SMA) at sub-arcsecond spatial resolution.
Results. The SMA resolved the hot molecular core into six submm continuum sources with the finest spatial resolution of 0.36′′ × 0.25′′
(∼1800 AU) achieved so far. Four of them located within 7800 (AU)2 comprise a proto-Trapezium system with estimated protostellar densities
of 1.4× 105 protostars/pc3 . The plethora of ∼ 80 spectral lines allows us to study the molecular outflow(s), the core kinematics, the temperature
structure of the region as well as chemical effects. The derived hot core temperatures are of the order 300 K. We find interesting chemical
spatial differentiations, e.g., C34S is deficient toward the hot core and is enhanced at the UCHii/ hot core interface, which may be explained
by temperature sensitive desorption from grains and following gas phase chemistry. The SiO(8–7) emission outlines likely two molecular
outflows emanating from this hot core region. Emission from most other molecules peaks centrally on the hot core and is not dominated by
any individual submm peak. Potential reasons for that are discussed. A few spectral lines that are associated with the main submm continuum
source, show a velocity gradient perpendicular to the large-scale outflow. Since this velocity structure comprises three of the central protostellar
sources, this is not a Keplerian disk. While the data are consistent with a gas core that may rotate and/or collapse, we cannot exclude the
outflow(s) and/or nearby expanding UCHii region as possible alternative causes of this velocity pattern.
Key words. stars: formation – ISM: jets and outflows – ISM: molecules – stars: early-type – stars: individual (G29.96−0.02) – (stars:) binaries
(including multiple): close
1. Introduction
Hot molecular cores represent an early evolutionary stage in
massive star formation prior to the formation of an ultra-
compact Hii region (UCHii). Single-dish line surveys toward
hot cores have revealed high abundances of many molecu-
lar species and temperatures usually exceeding 100 K (e.g.,
Schilke et al. 1997; Hatchell et al. 1998; McCutcheon et al.
2000). Unfortunately, most hot cores are relatively far away
(a few kpc, Orion-KL being an important exception), and high-
spatial resolution studies are important to disentangle the var-
Send offprint requests to: H. Beuther
ious components in the region, to resolve potential multiple
heating sources, and to search for chemical variations through-
out the regions. Here we present sub-arcsecond resolution
submm spectral line and dust continuum observations of the
hot core G29.96−0.02, characterizing the physical and chemi-
cal properties of this prototypical region.
The hot core/UCHii region G29.96−0.02 is a well studied
source comprising a cometary UCHii region and approximately
2.6′′ to the west a hot molecular core (Wood & Churchwell
1989; Cesaroni et al. 1994, 1998). G29.96−0.02 is at a dis-
tance of ∼6 kpc (Pratap et al. 1999), the bolometric luminos-
ity measured with IRAS is very high with L ∼ 1.4 × 106 L⊙
http://arxiv.org/abs/0704.0518v1
2 Beuther et al.: SMA observations of G29.96−0.02
(Cesaroni et al. 1994). Since the region harbors at least two
massive (proto)stars (within the UCHii region and the hot core)
this luminosity must be distributed over various sources. Based
on cm continuum free-free emission, Cesaroni et al. (1994) cal-
culate a luminosity for the UCHii region of Lcm ∼ 4.4×10
5 L⊙.
Furthermore, they try to estimate the luminosity of the hot core
via a first order black-body approximation and get a value of
Lbb ∼ 1.2×10
5 L⊙. Later, Olmi et al. (2003) derive a similar es-
timate (∼ 9 × 104 L⊙) via integrating a much better determined
SED. The exciting source of the UCHii region has been identi-
fied in the near-infrared as an O5-O8 star (Watson & Hanson
1997). Furthermore, Pratap et al. (1999) identified two addi-
tional sources toward the rim of the UCHii region and an en-
hanced density of reddened sources indicative of an embedded
cluster.
A line survey toward a number of UCHii regions reveals
that G29.96−0.02 is a strong molecular line emitter in nearly
all observed species (Hatchell et al. 1998). High-angular res-
olution studies show that many species (e.g., NH3, CH3CN,
HNCO, HCOOCH3) peak toward the main H2O maser cluster
∼ 2.6′′ west of the UCHii region (e.g, Hofner & Churchwell
1996; Cesaroni et al. 1998; Olmi et al. 2003), whereas CH3OH
peaks ∼ 4′′ further south-west associated with another iso-
lated H2O maser feature (Pratap et al. 1999). Hoffman et al.
(2003) detected one of the relatively rare H2CO masers toward
the hot core position. These masers are proposed to trace the
warm molecular gas in the vicinity of young forming massive
stars (Araya et al. 2006). The signature of a CH3OH peak off-
set from the other molecular lines is reminiscent of Orion-KL
(e.g., Wright et al. 1996; Beuther et al. 2005b). Temperature
estimates toward the hot core based on high-density trac-
ers vary between 80 and 150 K (e.g., Cesaroni et al. 1994;
Hatchell et al. 1998; Pratap et al. 1999; Olmi et al. 2003).
While Gibb et al. (2004) detect a molecular outflow in
H2S emanating from the hot core in approximately the south-
east north-west direction, Cesaroni et al. (1998) and Olmi et al.
(2003) detect a velocity gradient in the east-west direction
in the high-density tracers NH3(4,4) and CH3CN, consistent
with a rotating disk around an embedded protostar. However,
Maxia et al. (2001) also report that their rather low-resolution
5.9′′ × 3.7′′ (≈ 0.15 pc) SiO(2–1) data are consistent with the
disk scenario as well. This is a bit puzzling since SiO is usually
found to trace shocked gas in outflows and not more quies-
cent gas in disks. Inspecting their SiO image again (Fig. 6 in
Maxia et al. 2001), this interpretation is not unambiguous, the
data also appear to be consistent with the outflow observed in
H2S (Gibb et al. 2004). It is possible that the spatial resolution
of their SiO(2–1) observations is not sufficient to really disen-
tangle the outflow in this distant region.
Olmi et al. (2003) compiled the SED from cm to mid-
infrared wavelengths. While the 3 mm data are still strongly
dominated by the free-free emission (Olmi et al. 2003), at
1 mm the hot core becomes clearly distinguished from the ad-
jacent UCHii region (Wyrowski et al. 2002). G29.96−0.02 is
one of the few hot cores which is detected at mid-infrared
wavelengths (De Buizer et al. 2002). Interestingly, the mid-
infrared peak is ∼ 0.5′′ (∼3000 AU) offset from the NH3(4,4)
hot core position. While Gibb et al. (2004) speculate that the
mid-infrared peak might arise from the scattered light only,
De Buizer et al. (2002) suggest that it could trace a second mas-
sive source within the same core. This hypothesis can be tested
via very-high-angular-resolution submm continuum studies.
2. Observations
We have observed the hot core G29.96−0.02 with the
Submillimeter Array (SMA1, Ho et al. 2004) during four nights
between May and November 2005. We used all available ar-
ray configurations (compact, extended, very extended, for de-
tails see Table 1) with unprojected baselines between 16 and
500 m, resulting at 862 µm in a projected baseline range from
16.5 to 591 kλ. The chosen phase center was the peak position
of the associated UCHii region R.A. [J2000.0]: 18h46m03.s99
and Decl. [J2000.0] −02◦39′21.′′47. The velocity of rest is
vlsr ∼ +98 km s
−1 (Churchwell et al. 1990).
Table 1. Observing parameters
Date Config. # ant. Source loop τ(225GHz)
[hours]
28.May05 very ext. 6 7.0 0.13-0.16
18.Jul.05 comp. 7 7.5 0.06-0.09
4.Sep.05 ext. 6 4.5 0.06-0.08
5.Nov.05 very ext. 7 3.0 0.06
For bandpass calibration we used Ganymede in the com-
pact configuration and 3C279 and 3C454.3 in the extended and
very extended configuration. The flux scale was derived in the
compact configuration again from observations of Ganymede.
For two datasets of the more extended configurations, we used
3C454.3 for the relative scaling between the various baselines
and then scaled that absolutely via observations of Uranus.
For the fourth dataset we did the flux calibration using 3C279
only. The flux accuracy is estimated to be accurate within 20%.
Phase and amplitude calibration was done via frequent observa-
tions of the quasars 1743-038 and 1751+096, about 15.5◦ and
18.3◦ from the phase center of G29.96−0.02. The zenith opac-
ity τ(348GHz), measured with the NRAO tipping radiometer
located at the Caltech Submillimeter Observatory, varied dur-
ing the different observation nights between ∼0.15 and ∼0.4
(scaled from the 225 GHz measurement). The receiver operated
in a double-sideband mode with an IF band of 4-6 GHz so that
the upper and lower sideband were separated by 10 GHz. The
central frequencies of the upper and lower sideband were 348.2
and 338.2 GHz, respectively. The correlator had a bandwidth
of 2 GHz and the channel spacing was 0.8125 MHz. Measured
double-sideband system temperatures corrected to the top of
the atmosphere were between 110 and 800 K, depending on the
zenith opacity and the elevation of the source. Our sensitivity
was dynamic-range limited by the side-lobes of the strongest
1 The Submillimeter Array is a joint project between the
Smithsonian Astrophysical Observatory and the Academia Sinica
Institute of Astronomy and Astrophysics, and is funded by the
Smithsonian Institution and the Academia Sinica.
Beuther et al.: SMA observations of G29.96−0.02 3
emission peaks and thus varied between the line maps of dif-
ferent molecules and molecular transitions. This limitation was
mainly due to the incomplete sampling of short uv-spacings
and the presence of extended structures. The 1σ rms for the
velocity-integrated molecular line maps (the velocity ranges for
the integrations were chosen for each line separately depend-
ing on the line-widths and intensities) ranged between 36 and
76 mJy. The average synthesized beam of the spectral line maps
was 0.65′′×0.48′′ (P.A. −83◦). The 862 µm submm continuum
image was created by averaging the apparently line-free parts
of the upper sideband. The 1σ rms of the submm continuum
image was ∼ 21 mJy/beam, and the achieved synthesized beam
was 0.36′′×0.25′′ (P.A. 18◦), the smallest beam obtained so far
with the SMA. The different synthesized beams between line
and continuum maps are due to different applied weightings in
the imaging process (“robust” parameters set in MIRIAD to 0
and -2, respectively) because there was insufficient signal-to-
noise in the line data obtained in the very extended configura-
tion. The initial flagging and calibration was done with the IDL
superset MIR originally developed for the Owens Valley Radio
Observatory (Scoville et al. 1993) and adapted for the SMA2.
The imaging and data analysis were conducted in MIRIAD
(Sault et al. 1995).
3. Results
3.1. Submillimeter continuum emission
Figure 1 presents the 862 µm continuum emission extracted
from the line-free parts of the upper sideband spectrum
(∼1.8 GHz in total used) shown in Figure 4. The very high spa-
tial resolution of 0.36′′ × 0.25′′ corresponds to a linear res-
olution of ∼ 1800 AU at the given distance of ∼6 kpc. The
submm continuum emission peaks approximately 2′′ west of
the UCHii region and is associated with the molecular line
emission known from previous observations. We do not de-
tect any submm continuum emission toward the UCHii re-
gion itself. At the given spatial resolution, for the first time
multiplicity within the G29.96−0.02 hot core is resolved and
we identify 6 submm continuum emission peaks (submm1 to
submm6) above the 3σ level of 63 mJy beam−1 (Fig. 1). We
consider submm1 and submm2 to be separate sources instead
of a dust ridge because we count compact spherical or ellip-
tical sources and their emission peaks are separated by about
one synthesized beam. The four strongest submm peaks, that
are all > 6σ detections, are located within a region of (1.3′′)2
(7800 AU) in diameter. The submm peak submm1 is associated
with H2O and H2CO maser emission (Hofner & Churchwell
1996; Hoffman et al. 2003), and we consider this to be probably
the most luminous sub-source. The other H2O maser peaks are
offset from the submm continuum emission. The mid-infrared
source detected by De Buizer et al. (2002) is offset > 1′′ from
the submm emission. This may either be due to uncertainties
in the MIR astrometry or the MIR emission may trace another
young source in the region. It should be noted that the class ii
CH3OH masers detected by Minier et al. (2001) peak close to
2 The MIR cookbook by Charlie Qi can be found at
http://cfa-www.harvard.edu/∼cqi/mircook.html.
the MIR source as well, which indicates that the MIR offset
from the hot core may well be real.
Table 2 lists the absolute source positions, their 862µm
peak intensities and the integrated flux densities approximately
associated with each of the sub-sources. Calculating the bright-
ness temperature Tb of the corresponding Planck-function for,
e.g., submm1, we get Tb(Peak1) ∼ 27 K. Assuming hot core
dust temperatures of ∼ 100 K, the usual assumption of opti-
cally thin dust emission is not really valid anymore, and one
gets an approximate beam-averaged optical depth τ of the
dust emission of ∼0.3. To calculate the dust and gas masses,
we can follow the mass determination outlined in Hildebrand
(1983) and Beuther et al. (2002, 2005a), which assumes op-
tically thin emission, and correct that for the increased dust
opacity. Assuming constant emission along the line of sight,
the opacity correction factor C is
1 − e−τ
With τ ∼ 0.3, we get a correction factor C ∼ 1.16 still compa-
rably small. Assuming a dust opacity index β = 1.5, the dust
opacity per unit dust mass is κ(862µm) ∼ 1.5 cm2g−1 (with
the reference value κ(250µm) ∼ 9.4 cm2g−1, see Hildebrand
1983), and we assume a gas-to-dust ratio of 100. Given the
uncertainties in β and T , we estimate the masses to be accu-
rate within a factor 4. Table 2 gives the derived masses and
beam-averaged column densities. Each sub-peak has a mass
of a few M⊙, and the main submm1 exhibits approximately
10 M⊙ of compact, warm gas and dust emission. The inte-
grated 862 µm continuum flux density of the central region
comprising the four main submm continuum sources amounts
to 1.16 Jy. At an average dust temperature of 100 K, this cor-
responds to a central core mass of 39.9 M⊙. In comparison to
these flux density measurements, Thompson et al. (2006) ob-
served with SCUBA 850µm peak and integrated flux densities
of ∼ 11.5 ± 1.2 Jy/(14′′beam) and ∼19.2 Jy, respectively. The
ratio between peak and integrated JCMT fluxes already indi-
cates non-compact emission even on that scales. Furthermore,
subtracting a typical line contamination of the continuum emis-
sion in hot cores of the order 25% (e.g., NGC6334I, Hunter et
al. in prep.), the total 850µm single-dish continuum flux den-
sity should amount to ∼8.6 Jy. Compared with the integrated
flux density in the SMA data of ∼1.74 Jy, this indicates that
approximately 80% of the single-dish emission is filtered out
by the missing short spacings in the interferometer data. The
dust and gas in the central region have higher temperatures than
the components filtered out on larger spatial scales, and since
the dust and gas mass is inversely proportionally related to the
temperature by MH2 ∝ (e
hν/kT − 1) (e.g., Beuther et al. 2002),
a greater proportion of the mass (> 80%) is filtered out in the
SMA data. However, the SMA image reveals the most compact
hot gas and dust cores at the center of the evolving massive
star-forming region. The shortest baseline of the SMA obser-
vations of ∼16.5 kλ correspond to scales > 12′′ which hence
have to be filtered out entirely. However, even smaller scales are
missing because the uv-spacings corresponding to scales ≥ 5′′
are still relatively poorly sampled and the image presented in
Figure 1 is only sensitive to spatial scales of the order a few
http://cfa-www.harvard.edu/~cqi/mircook.html
4 Beuther et al.: SMA observations of G29.96−0.02
Fig. 1. The hot core UCHii region G29.96−0.02. The grey-scale with contours shows the submm continuum emission with a
spatial resolution of 0.36′′×0.25′′. The contour levels start at the 1σ level of 21 mJy beam−1 and continue at 63, 105 mJy beam−1
(black contours) to 147, 168 mJy beam−1 (white contours). The dashed contours outline the cm continuum emission from the
UCHii region and the thick contours show the NH3 emission (Cesaroni et al. 1994). The contouring is done from 15 to 95% (step
10%) of the peak emission of each image, respectively (S peak(1.2cm) = 109mJy/beam, S peak(NH3) = 15mJy/beam). Triangles,
circles and pentagons show the H2O (Hofner & Churchwell 1996), H2CO (Hoffman et al. 2003) and class ii CH3OH (Minier et al.
2001) maser positions. The star marks the peak of the MIR emission (De Buizer et al. 2002), which is not a point source but has
a similar size as the NH3 emission. The squares mark the infrared sources by Pratap et al. (1999).
arcseconds. The submm peaks detected by the SMA are much
stronger than what would have been expected if the single-dish
flux (∼8.6 Jy) were uniformly distributed over the SCUBA pri-
mary beam of 14′′, even ignoring any spatial filtering and miss-
ing flux effects (This would result in ∼ 4 mJy per synthesized
SMA beam.). This shows that the emission measured on the
small spatial scales sampled by the SMA represents the com-
pact core emission much better than expected. However, it does
not imply that the gas masses measured by the SMA are the
only gas reservoir the embedded protostars have for their on-
going accretion; they may also gain mass from the large-scale
gas envelope that is filtered out by our observations (see also
the competitive accretion scenario, e.g., Bonnell et al. 2004).
The derived beam-averaged H2 column densities are of the or-
der a few times 1024 cm−2, corresponding to visual extinctions
Av of a few 1000 (Av = NH/0.94 × 10
21, Frerking et al. 1982).
3.2. Spectral line emission
Figure 4 presents spectra extracted toward the main submm
submm1 with an angular resolution of 0.64′′×0.47′′ compared
to the submm continuum map (see §2). More than 80 spectral
lines from 18 molecular species, isotopologues or vibrationally
excited species have been identified with a minor fraction of
∼5% of unidentified lines (UL) (Tables 6 & 4). The range of up-
per level excitation temperatures for the many lines varies be-
tween approximately 40 and 750 K (Table 6). Therefore, with
one set of observations we are able to trace various gaseous
temperature components from the relatively colder gas sur-
Beuther et al.: SMA observations of G29.96−0.02 5
Table 2. Submm continuum source parameters
Source R.A. Dec. S peak S int M N
[J2000] [J2000] [ mJy
] [mJy] [M⊙] [10
24cm−3]
submm1 18:46:03.786 -02:39:22.19 173 288 11.5 5.7
submm2 18:46:03.789 -02:39:22.48 168 237 9.5 5.5
submm3 18:46:03.758 -02:39:22.16 138 178 7.1 4.5
submm4 18:46:03.736 -02:39:22.65 151 249 9.9 5.0
submm5 18:46:03.710 -02:39:23.33 68 106 4.2 2.2
submm6 18:46:03.665 -02:39:23.80 84 85 3.4 2.8
The Table shows the peak intensities S peak, the integrated intensities S int, the derived gas masses M as well as the H2 column densities N.
rounding the hot core region to the densest and warmest gas
best observed in some of the vibrationally excited lines.
Table 3. Peak intensities, rms and velocity ranges for images
in Figs. 2 & 3.
Line S peak rms ∆v
mJy/beam mJy/beam km/s
862µm cont., low res. 422 17
CH3OH(73,5 − 62,4) 878 64 [90,104]
13CH3OH(137,7 − 127,6) 752 51 [95,101]
CH3OH(74,3 − 64,3), vt = 1 1419 69 [91,105]
CH3OCH3(74,3 − 63,4) 669 46 [94,104]
C2H5OH(157,9 − 156,10) 586 51 [95,100]
SiO(8 − 7) 391 36 [75,105]
C34S(7 − 6) 592 62 [92,104]
H2CS(101,0 − 91,9) 933 69 [92,100]
34SO(88 − 77) 827 57 [95,103]
SO2(144,14 − 183,15) 544 53 [94,100]
HCOOCH3(275,22 − 265,21) 491 70 [96,100]
CH3CN(198 − 188) 788 71 [94,100]
CH3CH2CN(383,36 − 373,35) 791 56 [94,102]
CH3CHCN(362,34 − 352,32) 655 68 [96,100]
HC3N(37 − 36), v7 = 1 622 55 [94,102]
HC3N(37 − 36), v7 = 2 416 57 [94,100]
HN13C(4 − 3) 1149 76 [94,100]
Figures 2 and 3 now present integrated images of the var-
ious detected species, isotopologues and vibrationally excited
lines. For comparison, Figure 2 also shows the submm con-
tinuum emission reduced with the same degraded spatial res-
olution as the line images. All images show emission in the
vicinity of the hot molecular core and no emission toward the
associated UCHii region. However, the morphology varies sig-
nificantly between many of the observed molecular line maps.
The molecular emission is largely confined to the central region
of the main four submm continuum peaks, and we do not detect
appreciable molecular emission toward the continuum peaks
5 and 6. Reducing the submm continuum data with the same
spatial resolution as the line images, the four submm peaks
are smoothed to a single elongated structure peaking close to
the submm peak submm1 (Fig. 2, top-left panel). The ground
state CH3OH emission is relatively broadly distributed with
two peaks in east-west direction, and one may associate one
with the submm peaks 1 and 2 and the other with the submm
peak submm3, but most other maps show on average one spec-
tral line peak somewhere in the middle of the 4 main submm
continuum peaks, similar to the lower-resolution submm con-
tinuum map.
However, there are also a few species which significantly
deviate from this picture and show a different spatial morphol-
ogy. For example SiO is more extended in north-east south-
west direction likely due to a molecular outflow (§3.3). Also
interesting is the emission from C34S which lacks emission
around the central four submm peaks but is stronger in the in-
terface region between the hot molecular core and the UCHii
region (§4.3). Furthermore, there are a few spectral line maps
– mainly those from likely optically thin lines (HCOOCH3,
HN13C), highly excited lines (CH3CHCN) and vibrationally
excited lines (CH3OH vt = 1, 2, HC3N v7 = 1, 2) – which show
their emission peaks concentrated toward the main submm
peak submm1 (§4.5).
Previous lower-resolution (∼ 10′′) molecular line observa-
tions revealed strong CH3OH emission toward the H2O maser
feature approximately 4′′ south-west of the hot core peak
(Fig. 1, Hofner & Churchwell 1996; Pratap et al. 1999). A lit-
tle bit surprising, we do not detect any CH3OH emission (nor
any other species) toward that south-western position, even
when imaged at low angular resolution using only the com-
pact configuration data (therefore, we do not cover that posi-
tion in Figures 2 and 3). Pratap et al. (1999) discuss mainly
two possibilities to explain this discrepancy: Either their ob-
served specific CH3OH(80 − 71) line is a weak maser and
we do not cover any comparable CH3OH line, or the emis-
sion covered by the lower-resolution data is relatively extended
and filtered out by our observations. As discussed in the pre-
vious section, the shortest baseline of our observations was
∼16 m, implying that we are not sensitive to any scales > 12′′.
Since the CH3OH emission in Pratap et al. (1999) is slightly
resolved by their synthesized beam of 12.6′′ × 9.8′′, it is un-
likely that we would have filtered out all emission. However,
among the many observed CH3OH lines (Table 6), some have
similar excitation temperatures of the order 80 K as the line
observed by Pratap et al. (1999), and we would expect to de-
tect thermal emission from these lines as well. Therefore, our
non-detection of CH3OH emission toward the south-western
H2O maser position supports rather their suggested scenario of
weak CH3OH maser emission in the previously reported obser-
vations (Pratap et al. 1999).
6 Beuther et al.: SMA observations of G29.96−0.02
Fig. 4. Lower and upper sideband spectra extracted toward the
submm1. The spatial resolution of these data is 0.64′′ × 0.47′′.
The main line identifications are shown in both panels.
Table 4. Detected molecular species
Species Isotopologues Vibrational states
CH3OH
13CH3OH CH3OH, vt = 1, 2
CH3OCH3
C2H5OH
HCOOCH3
CH3CN
CH3CH2CN
CH3CHCN
HC3N, v7 = 1, 2
HN13C
a The detection of this CH3OH vt = 2 line is
doubtful since other close vt = 2 lines with
similar excitation temperatures were not detected.
3.3. Molecular outflow emission
The SiO(8-7) spectrum spans a large range of velocities from
∼75 to ∼111 km s−1. Integrating the blue- and red-shifted emis-
sion, one gets the outflow map presented in Fig. 5. The elon-
gated north-west south-east structure is consistent with the pre-
viously proposed outflow by Gibb et al. (2004). The additional
red feature north-east of the central hot core region makes the
interpretation ambiguous: If the north-west south-east outflow
is a relatively highly collimated jet, then the north-eastern red
feature could be attributed to an additional outflow leaving
the core in north-east south-west direction. The blue wing of
that potential second outflow would not detected in our data.
However, since we are filtering out any larger-scale emission,
it is also possible that the red SiO features south-east and north-
east of the main core are part of the same wide-angle outflow
tracing potentially the limb-brightened cavity walls. In this sce-
nario, our observations would miss part of the blue-shifted
wide-angle outflow lobe. With the current data, it is difficult to
clearly distinguish between the two scenarios. However, com-
paring the elongated blue-shifted SiO(8–7) data with the pre-
vious north-west south-eastern outflow observed in H2S by
Gibb et al. (2004), it appears that this is the most likely direc-
tion of the main outflow of the region. Therefore, the multi-
ple outflow scenario appears more likely for the hot core in
G29.96−0.02. The lower resolution SiO(2–1) observation by
Maxia et al. (2001) are also consistent with this scenario. Based
on these data, we cannot conclusively say which of the submm
continuum sources submm1 to submm4 contribute to driving
the outflows.
Fig. 5. SiO(8-7) outflow map. The full and dashed contours
are integrated over the blue- and redshifted SiO emission as
shown in the figure. The contouring starts at ±2σ and continues
in ±1σ steps (thick contours positive, thin contours negative).
The 1σ values for the blue- and red-shifted images are 48 and
46 mJy beam−1, respectively. The markers are the same as in
the previous images, the synthesized beam of 0.68′′ × 0.49′′ is
shown at the bottom right, and the arrows guide the eye for the
potential directions of the two discussed outflows. The offsets
on the axes are relative to the phase center.
4. Discussion
4.1. The formation of a proto-Trapezium system?
The four main submm continuum peaks are located within a
projected area of 7800 × 7800 (AU)2 on the sky. The projected
separation ∆θ between individual sub-sources varies between
1800 AU (peaks 1 and 2) and 5400 AU (peaks 1 and 4, see Table
5). Could the four central submm peaks be the predecessors
of a future Trapezium system? Trapezia are defined as non-
hierarchical multiple systems of three or more stars where the
Beuther et al.: SMA observations of G29.96−0.02 7
largest projected separation between Trapezia members should
not exceed the smallest projected separation by a factor of 3
(Sharpless 1954; Ambartsumian 1955; Abt & Corbally 2000).
This criterion is satisfied by the four submm peaks at the cen-
ter of the G29.96−0.02 hot core. The 14 optically identified
Trapezia discussed by Abt & Corbally (2000) have mean radii
to the furthest outlying member of ∼ 4 × 104 AU, with the
largest radius of ∼ 5.4 × 105 AU (∼2.6 pc), the approximate
dimension of an open cluster. Therefore, the protostellar pro-
jected separations of the tentative proto-Trapezium candidate in
G29.96−0.02 are significantly smaller than in typical optically
visible Trapezia systems. A similar small size for a candidate
Trapezium system has recently been reported for the multiple
system in W3IRS5 (Megeath et al. 2005).
Table 5. Spatial separation
Pair ∆θ ∆x
[′′] [AU]
1-2 0.3 1800
1-3 0.5 3000
1-4 0.9 5400
2-3 0.6 3600
2-4 0.8 4800
3-4 0.6 3600
The numbers in column 1 correspond to the numbers of the submm
peaks.
The small sizes of the proto-Trapezia in G29.96−0.02 and
W3IRS5 may be attributed to their youth. During their upcom-
ing evolution, these young system will expel most of the sur-
rounding gas and dust envelope via the protostellar outflows
and strong uv-radiation. Therefore, the whole gravitational po-
tential of the system will decrease and the kinetic energy may
dominate. Systems with positive total energy will globally ex-
pand and will eventually be observable as a larger-scale optical
Trapezia systems (Ambartsumian 1955).
With the given data it is hard to estimate how massive the
expected Trapezia stars are and will finally be at the end of
their formation processes. The integrated hot core luminosity
is estimated to be ∼ 105 L⊙ (Cesaroni et al. 1994; Olmi et al.
2003), in contrast to the integrated luminosity of the whole re-
gion measured by the large IRAS beam of ∼ 106 L⊙. Producing
105 L⊙ requires either an O7 star or a few stars of comparable
but lower masses. Nevertheless, the numbers imply that this
Trapezium system should form at least one or more massive
stars. Although the gas masses we derived from our dust con-
tinuum data (Table 2) are relatively low, that does not neces-
sarily imply that their mass reservoir is restricted to these gas
masses because it is possible that they may accrete additional
gas from the larger-scale envelope that is filtered out by our
observations. This scenario is predicted by the competitive ac-
cretion model for massive star formation (e.g., Bonnell et al.
2004). The fact that the gas masses we find for the four
strongest submm sources are all similar allows to speculate that
they may form about similar mass stars in the end, however, this
cannot be proven by these data in more detail.
Assuming that the projected size of the potential
proto-Trapezium system in G29.96−0.02 of approximately
7800 (AU)2 resembles a 3-dimensional sphere of radius
∼3900 AU, we can estimate the current protostellar volume
density of the region to approximately 1.4 × 105 protostars
per cubic pc. This number is larger than typical stellar den-
sities in young clusters of the order 104 stars per cubic pc
(Lada & Lada 2003), but it is still below the extremely high
(proto)stellar densities required for protostellar merger models
of the order 106 to 108 stars per cubic pc (Bonnell et al. 1998,
2004; Stahler et al. 2000; Bally & Zinnecker 2005).
Although we have not yet observed the extremely high
(proto)stellar densities predicted by the coalescence scenario,
as soon as we observe massive star-forming regions with a spa-
tial resolution ≤ 4000 AU, we begin to resolve multiplicity and
potential proto-Trapezia (see also the recent observations of
NGC6334I and I(N) by Hunter et al. 2006). Furthermore, this
(proto)stellar density may even be a lower limit, since we ob-
serve only a two-dimensional projection and are additionally
sensitivity limited to masses ≥ 2.1 M⊙ (corresponding to the
3σ flux limit of 63 mJy beam−1 at the assumed temperature of
100 K). Higher spatial resolution has so far always increased
the observed (proto)stellar densities, and it is possible that in
the future we may reach the 106 requirement for merging to
play a role. However, it is also important to get better theoret-
ical predictions of potential merger signatures that observers
could look for.
4.2. Various episodes of massive star formation?
It is interesting to note that the previously identified mid-
infrared source (De Buizer et al. 2002) is offset from the
submm continuum peaks. Although the mid-infrared astrome-
try is usually relatively uncertain, the association of the mid-
infrared peak with class ii CH3OH maser emission with an
absolute positional uncertainty of only 30 mas (Minier et al.
2001) is indicative that the offset may be real. Combining the
facts that we find within a small region of only ∼20000 AU
(∼0.1 pc) at least three different regions of massive star for-
mation – the UCHii region, the mid-infrared source, and the
submm continuum sources – indicates that not all massive stars
within the same evolving cluster are coeval but that sequences
of massive star formation may take place even on such small
spatial scales.
4.3. Carbon mono-sulfide C34S
One of the most striking spectral line maps is from the rare car-
bon mono-sulfide isotopologue C34S(7–6). Its emission peak is
not toward the hot core nor any of the submm continuum peaks,
but largely east of it in the interface region between the submm
continuum peaks and the UCHii region. Hence, one likes to
understand why the C34S emission is that weak toward the hot
core region and that strong at the hot core/UCHii region inter-
face.
CS usually desorbs from dust grains at moderate temper-
atures of a few 10 K, hence it should be observable relatively
8 Beuther et al.: SMA observations of G29.96−0.02
early in the evolution of a growing hot molecular core (e.g.,
Viti et al. 2004). From 100 K upwards H2O is released from
grains, then it forms OH molecules, and the OH can react
with S to SO and SO2 (e.g., Charnley 1997). Therefore, the
initial high CS abundances should decrease with time while
the SO and SO2 are expected to increase with time (e.g.,
Wakelam et al. (2005)). As shown in Figure 2, 34SO peaks to-
ward the hot core where the derived CH3OH temperatures ex-
ceed the H2O evaporation temperature (see §4.4 and Fig. 7b,
potentially validating this theoretical prediction. According to
such chemical models, the hot core G29.96−0.02 should have
a chemical age of at least a few times 104 years.
The strong C34S emission in the hot core/UCHii interface
region may be explained in the same framework. In the molec-
ular evolution scheme outlined above, one would expect low
C34S emission toward the hot core with a maybe symmetrical
increase further-out. In the case of the G29.96−0.02 hot core,
we have the decrease toward the center, but the emission rises
only toward the east, north and west with the strongest increase
in the eastern hot core/UCHii region interface. If one compares
the C34S morphology in Figure 2 with the temperature distri-
bution in Figure 7b, one finds the lowest CH3OH temperatures
right in the vicinity of the C34S emission peaks, adding further
support to the proposed chemical picture.
Extrapolating this scenario to other molecules, it indicates
that species which are destroyed by H2O, e.g., molecular ions
such as HCO+ or N2H
+ (e.g., Bergin et al. 1998), are no good
probes of the inner regions of hot molecular cores.
4.4. Temperature structure
Leurini et al. (2004, 2007) investigated the diagnostic proper-
ties of methanol over a range of physical parameters typical
of high-mass star-forming regions. They found that the ground
state lines of CH3OH are mainly tracers of the spatial density
of the gas, although at submillimeter wavelengths high k tran-
sitions are also sensitive to the kinetic temperature. However,
in hot, dense regions such as hot cores, the effects of infrared
pumping on the level populations due to the thermal heating of
the dust is not negligible, but mimic the effect of collisional ex-
citation. For the ground state line, Leurini et al. (2007) found
that it is virtually impossible to distinguish between IR pump-
ing and pumping by collisions, as both mechanisms equally
populate the vt = 0 levels. On the other hand, the vibrationally
or torsionally excited lines have very high critical densities
(1010–1011 cm−3) and high level energies (T ≥ 200 K). They
are difficult to be populated by collisions and trace the IR field
instead.
To study the physical conditions of the gas around the main
continuum peaks in G29.96–0.02, we analyzed only the emis-
sion coming from the vt = 1 lines, as their optical depth is lower
than for the ground state, and their emission is confined to the
gas around the dust condensations, while the vt = 0 transitions
are more extended and can be affected by problems of missing
flux. We first fitted the methanol emission of the vt = 1 lines
(see Fig. 6) towards the peak position, using the method de-
scribed by Leurini et al. (2004, 2007) that is based on an LVG
analysis and includes radiative pumping (Leurini et al. 2007).
The continuum emission derived in §3.1 was used in the calcu-
lations to solve the equations for the level populations. The two
main dust condensations submm1 and submm2 fall in the beam
of the line data; however, we assumed that the emission is com-
ing from only one component, which is more extended than our
beam, and derived a CH3OH column density averaged over the
beam of 4×1017 cm−2. The corresponding methanol abundance,
relative to H2 is of the order of 10
−7, typical of hot core sources.
Since the emission from the vt = 1 lines is optically thin for
this column density, and also at higher values, we consider this
approach valid. The temperature derived toward the line peak
is 340 K. This corresponds to our best fit model, but from a χ2
analysis we can only infer a low limit of ∼220 K for the temper-
ature of the gas. Since lines are optically thin, the degeneracy
between kinetic temperature and column density is not solved,
and the model delivers good fit to the vt = 1 lines for lower
or higher temperatures by adjusting the methanol column den-
sity. However, the low temperature solutions (Tkin=100–200 K)
need high methanol abundances relative to H2(∼ 10
−6), which
can be hardly found at these temperatures. Moreover, lines are
optically thick for these column densities, and the assumption
of our analysis is not valid anymore.
We also investigated the line ratio between several vt = 1
lines at the column density derived for the main position, to
find the best temperature diagnostic tool among the methanol
lines and derive a temperature map of the region. We found
that the line ratios with the blend of lines at ∼ 337.64 GHz
increase with the temperature of the gas (Fig. 7a). However,
the blending of several transitions together complicates the use
of such diagnostic. In Fig. 7b, we show the map of the line ratio
between the 71,6 → 61,5-E vt = 1 at 337.708 GHz and the blend
between the 71,7 → 61,6-E vt = 1 at 337.642 GHz and 70,7 →
60,6-E vt = 1 at 337.644 GHz. Since line intensities do not
simply add up, we did not correct for the overlapping between
the two transitions. Two other lines, the 74,3 → 64,2-E vt = 1
and the 75,3 → 64,2-E vt = 1, are also very close in frequency.
This is seen in the linewidth of the blending, which is wider
than for the other lines. Therefore, we considered only half of
the channels of the blending at 337.64 GHz in our line ratio
analysis. From the ratio-map in Fig. 7b, submm1, submm2 and
submm3 of Table 1 show high temperatures (T≥ 300 K), while
relatively low temperature gas (T∼ 100 K) is found at R.A.
[J2000]=18h46m03s.818 Dec. [J2000]= −02◦39′22′′.14, close
to a secondary peak of many ground state lines of methanol
(Fig. 2). The temperature then decreases towards submm4. The
increase in the line ratio towards the south-east and north is
probably not true, but due to the poor signal to noise ratio in
these areas. Changes in the column densities along the area may
affect our results.
Fig. 6. Spectrum of the 7ka ,kb → 6ka ,kb−1 vt = 1 methanol band to-
ward the main dust condensation. Overlaid in black is the synthetic
spectrum resulting from the fit.
Beuther et al.: SMA observations of G29.96−0.02 9
Fig. 7. a: Modeled line ratio between the 71,6 → 61,5-E vt = 1 line and
the 71,7 → 61,6-E vt = 1 transitions, as function of the temperature. b:
Map of the line ratio between the same transitions in the inner region
around the peaks. The white stars mark the positions of the dust peaks;
the white dashed contours show the values of the line ratio from ∼ 150
to ∼ 350 K, which correspond to levels from 0.3 to 0.7 in step of 0.1
in the map. The solid black contours show the continuum emission
smoothed to the resolution of the line data (from 0.2 to 0.4 Jy/beam in
step of 0.05). The offsets on the axes are relative to the phase center.
4.5. Tracing rotation toward the massive cores
At the given lower spatial resolution of the spectral line data
compared to the submm continuum, we cannot resolve the four
submm peaks well. However, one of the aims of such multi-line
studies is to identify spectral lines that trace the massive proto-
stars and that are potentially associated with massive disk-like
structures. Such lines may then be used for kinematic gas stud-
ies of rotating gas envelopes, tori or accretion disks. Therefore,
we analyzed the data-cubes searching for velocity structures
indicative of any kind of rotation. In the large majority of spec-
tral lines, this was not successful and we could mostly not
identify coherent velocity structure. While chemical and tem-
perature effects (§4.3 & 4.4) may be responsible for parts of
that, the large column densities derived in §3.1 imply also large
molecular line column densities and hence large optical depths.
Therefore, many of the observed lines are likely optically thick
tracing only outer gas layers of the hot molecular core not pen-
etrating down to the deeply embedded protostars. Furthermore,
many molecules would not only be excited in the central ro-
tating disk-like structures but also in the surrounding envelope
and maybe the outflow. Hence, disentangling the different com-
ponents observationally remains a challenging task.
Fig. 8. Moment 1 maps of HN13C(4–3) (top) and HC3N(37–
36)v7 = 1 (bottom). The markers are the same as in the previous
images, and the synthesized beam of 0.68′′ × 0.49′′ is shown at
the bottom left. The offsets on the axes are relative to the phase
center.
The major exceptions are the molecular lines of the rare
isotopologue of hydrogen isocyanide HN13C(4–3) with a low
excitation temperature of only 42 K, and the vibrationally ex-
cited line of cyanoacetylene HC3N(37–36)v7 = 1 with a higher
excitation temperature of 629 K (Fig. 8). In both cases we find
a velocity gradient across the main submm peak submm1 with
a position angle of ∼ 45◦ from north. This is approximately
perpendicular to the molecular outflow discussed in §3.3 and
by Gibb et al. (2004). Interestingly, Gibb et al. (2004) also find
a similar velocity gradient in their central velocity channels of
H2S. The previously reported NH3 and CH3CN velocity gradi-
ents in approximately east-west direction (Cesaroni et al. 1998;
Olmi et al. 2003) have been observed with slightly lower spa-
tial resolution and are consistent with our data as well.
Our observations as well as previous work in the liter-
ature suggest that the G29.96−0.02 hot core exhibits a ve-
locity gradient in the dense gas in approximately north-east
south-west direction perpendicular to the molecular outflow
observed at larger scales. Based on the HN13C(4–3) map, the
diameter of this structure is ∼ 1.6′′ corresponding to radius
of ∼4800 AU. Since this emission encompasses not only the
submm peak submm1 but also the submm2 and submm3,
it is not genuine protostellar disk as often observed in low-
mass star-forming regions. The velocity structure does not re-
semble Keplerian rotation and may hence be due to some
larger-scale rotating envelope or torus that could transform
into a genuine accretion disks at smaller still unresolved spa-
tial scales (Cesaroni et al. 2007). Additional options to ex-
plain such a velocity gradient may be (a) interaction with the
2nd outflow in north-east–south-western direction, (b) inter-
action with the expanding UCHii region, and (c) global col-
lapse like recently proposed for NGC2264 (Peretto et al. 2006).
While we cannot exclude (a) and (b), option (c) of a globally
collapsing core appears particularly interesting because com-
bining rotation and collapse would result in an inward spi-
raling kinematic structure, potentially similar to the models
originally proposed for rotating low-mass cores (e.g., Ulrich
1976; Terebey et al. 1984). Recent hydrodynamic simulations
by Dobbs et al. (2005) and Krumholz et al. (2006) as well as
analytic studies by Kratter & Matzner (2006) find fragmenta-
tion and star formation within the massive disks forming early
in the collapse process of high-mass cores. This would be con-
sistent with the found three sub-sources (submm1 to submm3)
within the HN13C/HC3N structure. However, on a cautionary
note it needs to be stressed that the collapse/rotation scenario
is far from conclusive, and that the outflow and/or UCHii re-
gion can potentially influence the observed velocity pattern as
well. It remains puzzling that only these two lines exhibit the
discussed signatures whereas all the other spectral lines in our
setup do not.
5. Conclusions and Summary
The new 862 µm submm continuum and spectral line data ob-
tained with the SMA toward G29.96−0.02 at sub-arcsecond
spatial resolution resolve the hot molecular core into sev-
eral sub-sources. At an angular resolution of 0.36′′ × 0.25′′,
corresponding to linear scales of ∼1800 AU, the central core
contains four submm continuum peaks which resemble a
Trapezium-like multiple system at a very early evolutionary
stage. Assuming spherical symmetry for the hot core region,
the protostellar densities are high of the order 1.4 × 105 pro-
tostars per pc3. However, these protostellar densities are still
below the required values between 106 to 108 protostars/pc3 to
make coalescence of protostars a feasible process. Derived H2
column densities of the order a few 1024 cm−2 imply visual ex-
tinctions of a few 1000. The existence of three sites of massive
star formation in different evolutionary stages within a small
region (the UCHii region, the mid-infrared source, and the
submm continuum sources) indicates that sequences of mas-
sive star formation may take place within the same evolving
massive protocluster.
The 4 GHz of observed bandpass reveal a plethora of ap-
proximately 80 spectral lines from 18 molecular species, iso-
topologues or vibrationally excited lines. Only about 5% of the
10 Beuther et al.: SMA observations of G29.96−0.02
spectral lines remain unidentified. Most spectral lines peak to-
ward the hot molecular core, while a few species also show
more extended emission, likely due to molecular outflows and
chemical differentiation. The CH3OH line forest allows us to
investigate the temperature structure in more detail. We find
hot core temperatures≥ 300 K and decreasing temperature gra-
dients to the core edges. The SiO(8-7) observations confirm
a previously reported outflow Gibb et al. (2004) in north-west
south-east direction with a potential identification of a second
outflow emanating approximately in perpendicular direction.
Furthermore, C34S exhibits a peculiar morphology being weak
toward the hot molecular core and strong in its surroundings,
particular in the UCHii/hot core interface region. The C34S de-
ficiency toward the hot molecular core may be explained by
time-dependent chemical desorption from grains, where the
C34S desorbs early, and later-on after H2O desorbs from grains
forming OH, the sulphur reacts with the OH to form SO and
Furthermore, we were interested in identifying the best
molecular line tracers to investigate the kinematics and po-
tential disk-like structures in such dense and young massive
star-forming regions. Most spectral lines do not exhibit any
coherent velocity structure. A likely explanation for this un-
correlation between molecular line peaks and submm contin-
uum peaks is that many spectral lines may be optically thick in
such high-column-density regions, and that additional chem-
ical evolution and temperature effects complicate the picture.
Furthermore, many molecules are excited in various gas com-
ponents (e.g., disk, envelope, outflow), and it is often observa-
tionally difficult to disentangle the different contributions prop-
erly. There are a few exceptions of optically thin and vibra-
tionally excited lines that apparently probe deeper into the core
tracing submm1 better than other transitions. Investigating the
velocity pattern of these spectral lines, we find for some of
them a velocity gradient in the north-east south-west direc-
tion perpendicular to the molecular outflow. Since the spatial
scale of this structure is relatively large (∼4800 AU) compris-
ing three of the central protostellar sources, and since the veloc-
ity structure is not Keplerian, this is not a genuine Keplerian ac-
cretion disk. While these data are consistent with a larger-scale
toroid or envelope that may rotate and/or globally collapse, we
cannot exclude other explanations, such as that the influence
of the outflow(s) and/or expanding UCHII region produces the
observed velocity pattern. In addition to this, these data con-
firm previous findings that the high column densities, the large
optical depths of the spectral lines, the chemical evolution, and
the different spectral line contributions from various gas com-
ponents make it very difficult to identify suitable massive ac-
cretion disk tracers, and hence to study this phenomenon in a
more statistical fashion. (e.g., Beuther et al. 2006)
Acknowledgements. We like to thank Peter Schilke and Sebastian
Wolf for many interesting discussions about related subjects. We
also thank the anonymous referee whose comments helped improv-
ing the paper. H.B. acknowledges financial support by the Emmy-
Noether-Program of the Deutsche Forschungsgemeinschaft (DFG,
grant BE2578).
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12 Beuther et al.: SMA observations of G29.96−0.02
Fig. 2. Compilation of integrated line images (and submm continuum at the same spatial resolution) always shown in grey-scale
with contours and labeled at the bottom of each panel. The dashed contours show negative features due to missing short spacings.
The contouring is done from ±15 to ±95% (step ±10%) of the peak emission of each image, respectively. Peak fluxes S peak, rms
and integrated velocity ranges for each image are given in Table 3. The dotted contours again show the UCHii region and the
stars mark the submm continuum peaks from Figure 1. The offsets on the axes are relative to the phase center.
Beuther et al.: SMA observations of G29.96−0.02 13
Fig. 3. Continued Figure 2.
14 Beuther et al.: SMA observations of G29.96−0.02
Table 6. Line parameters
Freq. Line Eu Freq. Line Eu
GHz K GHz K
337.279 CH3OH(72,5 − 62,4)E(vt=2)
a 727 338.409 CH3OH(70,7 − 60,6)A 65
337.297 CH3OH(71,7 − 61,6)A(vt=1) 390 338.431 CH3OH(76,1 − 66,0)E 254
337.348 CH3CH2CN(383,36 − 373,35) 328 338.442 CH3OH(76,1 − 66,0)A 259
337.397 C34S(7–6) 65 CH3OH(76,2 − 66,1)A
− 259
337.421 CH3OCH3(212,19 − 203,18) 220 338.457 CH3OH(75,2 − 65,1)E 189
337.446 CH3CH2CN(374,33 − 364,32) 322 338.475 CH3OH(75,3 − 65,2)E 201
337.464 CH3OH(76,1 − 60,0)A(vt=1) 533 338.486 CH3OH(75,3 − 65,2)A 203
337.474 UL CH3OH(75,2 − 65,1)A
− 203
337.490 HCOOCH3(278,20 − 268,19)E 267 338.504 CH3OH(74,4 − 64,3)E 153
337.519 CH3OH(75,2 − 65,2)E(vt=1) 482 338.513 CH3OH(74,4 − 64,3)A
− 145
337.546 CH3OH(75,3 − 65,2)A(vt=1) 485 CH3OH(74,3 − 64,2)A 145
CH3OH(75,2 − 65,1)A
−(vt=1) 485 CH3OH(72,6 − 62,5)A
− 103
337.582 34SO(88 − 77) 86 338.530 CH3OH(74,3 − 64,2)E 161
337.605 CH3OH(72,5 − 62,4)E(vt=1) 429 338.541 CH3OH(73,5 − 63,4)A
+ 115
337.611 CH3OH(76,1 − 66,0)E(vt=1) 657 338.543 CH3OH(73,4 − 63,3)A
− 115
CH3OH(73,4 − 63,3)E(vt=1) 388 338.560 CH3OH(73,5 − 63,4)E 128
337.626 CH3OH(72,5 − 62,4)A(vt=1) 364 338.583 CH3OH(73,4 − 63,3)E 113
337.636 CH3OH(72,6 − 62,5)A
−(vt=1) 364 338.612 SO2(201,19 − 192,18) 199
337.642 CH3OH(71,7 − 61,6)E(vt=1) 356 338.615 CH3OH(71,6 − 61,5)E 86
337.644 CH3OH(70,7 − 60,6)E(vt=1) 365 338.640 CH3OH(72,5 − 62,4)A 103
337.646 CH3OH(74,3 − 64,2)E(vt=1) 470 338.722 CH3OH(72,5 − 62,4)E 87
337.648 CH3OH(75,3 − 65,2)E(vt=1) 611 338.723 CH3OH(72,6 − 62,5)E 91
337.655 CH3OH(73,5 − 63,4)A(vt=1) 461 338.760
13CH3OH(137,7 − 127,6)A 206
CH3OH(73,4 − 63,3)A
−(vt=1) 461 338.769 HC3N(37 − 36)v7 = 2 525
337.671 CH3OH(72,6 − 62,5)E(vt=1) 465 338.886 C2H5OH(157,8 − 156,19) 162
337.686 CH3OH(74,3 − 64,2)A(vt=1) 546 339.058 C2H5OH(147,7 − 146,8) 150
CH3OH(74,4 − 64,3)A
−(vt=1) 546 347.232 CH2CHCN(381,38 − 371,37) 329
CH3OH(75,2 − 65,1)E(vt=1) 494 347.331
28SiO(8–7) 75
337.708 CH3OH(71,6 − 61,5)E(vt=1) 489 347.446 UL
337.722 CH3OCH3(74,4 − 63,3)EE 48 347.494 HCOOCH3(275,22 − 265,21)A 247
337.732 CH3OCH3(74,3 − 63,3)EE 48 347.759 CH2CHCN(362,34 − 352,32) 317
337.749 CH3OH(70,7 − 60,6)A(vt=1) 489 347.792 UL
337.778 CH3OCH3(74,4 − 63,4)EE 48 347.842 UL
337.787 CH3OCH3(74,3 − 63,4)AA 48 347.916 C2H5OH(204,17 − 194,16) 251
337.825 HC3N(37 − 36)v7 = 1 629 347.983 UL
337.838 CH3OH(206,14 − 215,16)E 676 348.261 CH3CH2CN(392,37 − 382,36) 344
337.878 CH3OH(71,6 − 61,5)A(vt=2) 748 348.340 HN
13C(4–3) 42
337.969 CH3OH(71,6 − 61,5)A(vt=1) 390 348.345 CH3CH2CN(402,39 − 392,38) 351
338.081 H2CS(101,10 − 91,9) 102 348.532 H2CS(101,9 − 91,8) 105
338.125 CH3OH(70,7 − 60,6)E 78 348.910 HCOOCH3(289,20 − 279,19)E 295
338.143 CH3CH2CN(373,34 − 363,33) 317 348.911 CH3CN(199 − 189) 745
338.306 SO2(144,14 − 183,15) 197 349.025 CH3CN(198 − 188) 624
338.345 CH3OH(71,7 − 61,6)E 71 349.107 CH3OH(141,13 − 140,14) 43
338.405 CH3OH(76,2 − 66,1)E 244
a The detection of this CH3OH vt = 2 line is doubtful since other close vt = 2 lines with similar
excitation temperatures were not detected.
This figure "ch3oh1.jpg" is available in "jpg"
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This figure "hc3n_v1_mom1.jpg" is available in "jpg"
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http://arxiv.org/ps/0704.0518v1
	Introduction
	Observations
	Results
	Submillimeter continuum emission
	Spectral line emission
	Molecular outflow emission
	Discussion
	The formation of a proto-Trapezium system?
	Various episodes of massive star formation?
	Carbon mono-sulfide C34S
	Temperature structure
	Tracing rotation toward the massive cores
	Conclusions and Summary

  Aiming at a better understand of the physical and chemical processes in the
hot molecular core stage of high-mass star formation, we observed the
prototypical hot core G29.96-0.02 in the 862mu band with the Submillimeter
Array (SMA) at sub-arcsecond spatial resolution. The observations resolved the
hot molecular core into six submm continuum sources with the finest spatial
resolution of 0.36''x0.25'' (~1800AU) achieved so far. Four of them located
within 7800(AU)^2 comprise a proto-Trapezium system with estimated protostellar
densities of 1.4x0^5 protostars/pc^3. The plethora of ~80 spectral lines allows
us to study the molecular outflow(s), the core kinematics, the temperature
structure of the region as well as chemical effects. The derived hot core
temperatures are of the order 300K. We find interesting chemical spatial
differentiations, e.g., C34S is deficient toward the hot core and is enhanced
at the UCHII/hot core interface, which may be explained by temperature
sensitive desorption from grains and following gas phase chemistry. The
SiO(8-7) emission outlines likely two molecular outflows emanating from this
hot core region. Emission from most other molecules peaks centrally on the hot
core and is not dominated by any individual submm peak. Potential reasons for
that are discussed. A few spectral lines that are associated with the main
submm continuum source, show a velocity gradient perpendicular to the
large-scale outflow. Since this velocity structure comprises three of the
central protostellar sources, this is not a Keplerian disk. While the data are
consistent with a gas core that may rotate and/or collapse, we cannot exclude
the outflow(s) and/or nearby expanding UCHII region as possible alternative
causes of this velocity pattern.

Introduction
Hot molecular cores represent an early evolutionary stage in
massive star formation prior to the formation of an ultra-
compact Hii region (UCHii). Single-dish line surveys toward
hot cores have revealed high abundances of many molecu-
lar species and temperatures usually exceeding 100 K (e.g.,
Schilke et al. 1997; Hatchell et al. 1998; McCutcheon et al.
2000). Unfortunately, most hot cores are relatively far away
(a few kpc, Orion-KL being an important exception), and high-
spatial resolution studies are important to disentangle the var-
Send offprint requests to: H. Beuther
ious components in the region, to resolve potential multiple
heating sources, and to search for chemical variations through-
out the regions. Here we present sub-arcsecond resolution
submm spectral line and dust continuum observations of the
hot core G29.96−0.02, characterizing the physical and chemi-
cal properties of this prototypical region.
The hot core/UCHii region G29.96−0.02 is a well studied
source comprising a cometary UCHii region and approximately
2.6′′ to the west a hot molecular core (Wood & Churchwell
1989; Cesaroni et al. 1994, 1998). G29.96−0.02 is at a dis-
tance of ∼6 kpc (Pratap et al. 1999), the bolometric luminos-
ity measured with IRAS is very high with L ∼ 1.4 × 106 L⊙
http://arxiv.org/abs/0704.0518v1
2 Beuther et al.: SMA observations of G29.96−0.02
(Cesaroni et al. 1994). Since the region harbors at least two
massive (proto)stars (within the UCHii region and the hot core)
this luminosity must be distributed over various sources. Based
on cm continuum free-free emission, Cesaroni et al. (1994) cal-
culate a luminosity for the UCHii region of Lcm ∼ 4.4×10
5 L⊙.
Furthermore, they try to estimate the luminosity of the hot core
via a first order black-body approximation and get a value of
Lbb ∼ 1.2×10
5 L⊙. Later, Olmi et al. (2003) derive a similar es-
timate (∼ 9 × 104 L⊙) via integrating a much better determined
SED. The exciting source of the UCHii region has been identi-
fied in the near-infrared as an O5-O8 star (Watson & Hanson
1997). Furthermore, Pratap et al. (1999) identified two addi-
tional sources toward the rim of the UCHii region and an en-
hanced density of reddened sources indicative of an embedded
cluster.
A line survey toward a number of UCHii regions reveals
that G29.96−0.02 is a strong molecular line emitter in nearly
all observed species (Hatchell et al. 1998). High-angular res-
olution studies show that many species (e.g., NH3, CH3CN,
HNCO, HCOOCH3) peak toward the main H2O maser cluster
∼ 2.6′′ west of the UCHii region (e.g, Hofner & Churchwell
1996; Cesaroni et al. 1998; Olmi et al. 2003), whereas CH3OH
peaks ∼ 4′′ further south-west associated with another iso-
lated H2O maser feature (Pratap et al. 1999). Hoffman et al.
(2003) detected one of the relatively rare H2CO masers toward
the hot core position. These masers are proposed to trace the
warm molecular gas in the vicinity of young forming massive
stars (Araya et al. 2006). The signature of a CH3OH peak off-
set from the other molecular lines is reminiscent of Orion-KL
(e.g., Wright et al. 1996; Beuther et al. 2005b). Temperature
estimates toward the hot core based on high-density trac-
ers vary between 80 and 150 K (e.g., Cesaroni et al. 1994;
Hatchell et al. 1998; Pratap et al. 1999; Olmi et al. 2003).
While Gibb et al. (2004) detect a molecular outflow in
H2S emanating from the hot core in approximately the south-
east north-west direction, Cesaroni et al. (1998) and Olmi et al.
(2003) detect a velocity gradient in the east-west direction
in the high-density tracers NH3(4,4) and CH3CN, consistent
with a rotating disk around an embedded protostar. However,
Maxia et al. (2001) also report that their rather low-resolution
5.9′′ × 3.7′′ (≈ 0.15 pc) SiO(2–1) data are consistent with the
disk scenario as well. This is a bit puzzling since SiO is usually
found to trace shocked gas in outflows and not more quies-
cent gas in disks. Inspecting their SiO image again (Fig. 6 in
Maxia et al. 2001), this interpretation is not unambiguous, the
data also appear to be consistent with the outflow observed in
H2S (Gibb et al. 2004). It is possible that the spatial resolution
of their SiO(2–1) observations is not sufficient to really disen-
tangle the outflow in this distant region.
Olmi et al. (2003) compiled the SED from cm to mid-
infrared wavelengths. While the 3 mm data are still strongly
dominated by the free-free emission (Olmi et al. 2003), at
1 mm the hot core becomes clearly distinguished from the ad-
jacent UCHii region (Wyrowski et al. 2002). G29.96−0.02 is
one of the few hot cores which is detected at mid-infrared
wavelengths (De Buizer et al. 2002). Interestingly, the mid-
infrared peak is ∼ 0.5′′ (∼3000 AU) offset from the NH3(4,4)
hot core position. While Gibb et al. (2004) speculate that the
mid-infrared peak might arise from the scattered light only,
De Buizer et al. (2002) suggest that it could trace a second mas-
sive source within the same core. This hypothesis can be tested
via very-high-angular-resolution submm continuum studies.
2. Observations
We have observed the hot core G29.96−0.02 with the
Submillimeter Array (SMA1, Ho et al. 2004) during four nights
between May and November 2005. We used all available ar-
ray configurations (compact, extended, very extended, for de-
tails see Table 1) with unprojected baselines between 16 and
500 m, resulting at 862 µm in a projected baseline range from
16.5 to 591 kλ. The chosen phase center was the peak position
of the associated UCHii region R.A. [J2000.0]: 18h46m03.s99
and Decl. [J2000.0] −02◦39′21.′′47. The velocity of rest is
vlsr ∼ +98 km s
−1 (Churchwell et al. 1990).
Table 1. Observing parameters
Date Config. # ant. Source loop τ(225GHz)
[hours]
28.May05 very ext. 6 7.0 0.13-0.16
18.Jul.05 comp. 7 7.5 0.06-0.09
4.Sep.05 ext. 6 4.5 0.06-0.08
5.Nov.05 very ext. 7 3.0 0.06
For bandpass calibration we used Ganymede in the com-
pact configuration and 3C279 and 3C454.3 in the extended and
very extended configuration. The flux scale was derived in the
compact configuration again from observations of Ganymede.
For two datasets of the more extended configurations, we used
3C454.3 for the relative scaling between the various baselines
and then scaled that absolutely via observations of Uranus.
For the fourth dataset we did the flux calibration using 3C279
only. The flux accuracy is estimated to be accurate within 20%.
Phase and amplitude calibration was done via frequent observa-
tions of the quasars 1743-038 and 1751+096, about 15.5◦ and
18.3◦ from the phase center of G29.96−0.02. The zenith opac-
ity τ(348GHz), measured with the NRAO tipping radiometer
located at the Caltech Submillimeter Observatory, varied dur-
ing the different observation nights between ∼0.15 and ∼0.4
(scaled from the 225 GHz measurement). The receiver operated
in a double-sideband mode with an IF band of 4-6 GHz so that
the upper and lower sideband were separated by 10 GHz. The
central frequencies of the upper and lower sideband were 348.2
and 338.2 GHz, respectively. The correlator had a bandwidth
of 2 GHz and the channel spacing was 0.8125 MHz. Measured
double-sideband system temperatures corrected to the top of
the atmosphere were between 110 and 800 K, depending on the
zenith opacity and the elevation of the source. Our sensitivity
was dynamic-range limited by the side-lobes of the strongest
1 The Submillimeter Array is a joint project between the
Smithsonian Astrophysical Observatory and the Academia Sinica
Institute of Astronomy and Astrophysics, and is funded by the
Smithsonian Institution and the Academia Sinica.
Beuther et al.: SMA observations of G29.96−0.02 3
emission peaks and thus varied between the line maps of dif-
ferent molecules and molecular transitions. This limitation was
mainly due to the incomplete sampling of short uv-spacings
and the presence of extended structures. The 1σ rms for the
velocity-integrated molecular line maps (the velocity ranges for
the integrations were chosen for each line separately depend-
ing on the line-widths and intensities) ranged between 36 and
76 mJy. The average synthesized beam of the spectral line maps
was 0.65′′×0.48′′ (P.A. −83◦). The 862 µm submm continuum
image was created by averaging the apparently line-free parts
of the upper sideband. The 1σ rms of the submm continuum
image was ∼ 21 mJy/beam, and the achieved synthesized beam
was 0.36′′×0.25′′ (P.A. 18◦), the smallest beam obtained so far
with the SMA. The different synthesized beams between line
and continuum maps are due to different applied weightings in
the imaging process (“robust” parameters set in MIRIAD to 0
and -2, respectively) because there was insufficient signal-to-
noise in the line data obtained in the very extended configura-
tion. The initial flagging and calibration was done with the IDL
superset MIR originally developed for the Owens Valley Radio
Observatory (Scoville et al. 1993) and adapted for the SMA2.
The imaging and data analysis were conducted in MIRIAD
(Sault et al. 1995).
3. Results
3.1. Submillimeter continuum emission
Figure 1 presents the 862 µm continuum emission extracted
from the line-free parts of the upper sideband spectrum
(∼1.8 GHz in total used) shown in Figure 4. The very high spa-
tial resolution of 0.36′′ × 0.25′′ corresponds to a linear res-
olution of ∼ 1800 AU at the given distance of ∼6 kpc. The
submm continuum emission peaks approximately 2′′ west of
the UCHii region and is associated with the molecular line
emission known from previous observations. We do not de-
tect any submm continuum emission toward the UCHii re-
gion itself. At the given spatial resolution, for the first time
multiplicity within the G29.96−0.02 hot core is resolved and
we identify 6 submm continuum emission peaks (submm1 to
submm6) above the 3σ level of 63 mJy beam−1 (Fig. 1). We
consider submm1 and submm2 to be separate sources instead
of a dust ridge because we count compact spherical or ellip-
tical sources and their emission peaks are separated by about
one synthesized beam. The four strongest submm peaks, that
are all > 6σ detections, are located within a region of (1.3′′)2
(7800 AU) in diameter. The submm peak submm1 is associated
with H2O and H2CO maser emission (Hofner & Churchwell
1996; Hoffman et al. 2003), and we consider this to be probably
the most luminous sub-source. The other H2O maser peaks are
offset from the submm continuum emission. The mid-infrared
source detected by De Buizer et al. (2002) is offset > 1′′ from
the submm emission. This may either be due to uncertainties
in the MIR astrometry or the MIR emission may trace another
young source in the region. It should be noted that the class ii
CH3OH masers detected by Minier et al. (2001) peak close to
2 The MIR cookbook by Charlie Qi can be found at
http://cfa-www.harvard.edu/∼cqi/mircook.html.
the MIR source as well, which indicates that the MIR offset
from the hot core may well be real.
Table 2 lists the absolute source positions, their 862µm
peak intensities and the integrated flux densities approximately
associated with each of the sub-sources. Calculating the bright-
ness temperature Tb of the corresponding Planck-function for,
e.g., submm1, we get Tb(Peak1) ∼ 27 K. Assuming hot core
dust temperatures of ∼ 100 K, the usual assumption of opti-
cally thin dust emission is not really valid anymore, and one
gets an approximate beam-averaged optical depth τ of the
dust emission of ∼0.3. To calculate the dust and gas masses,
we can follow the mass determination outlined in Hildebrand
(1983) and Beuther et al. (2002, 2005a), which assumes op-
tically thin emission, and correct that for the increased dust
opacity. Assuming constant emission along the line of sight,
the opacity correction factor C is
1 − e−τ
With τ ∼ 0.3, we get a correction factor C ∼ 1.16 still compa-
rably small. Assuming a dust opacity index β = 1.5, the dust
opacity per unit dust mass is κ(862µm) ∼ 1.5 cm2g−1 (with
the reference value κ(250µm) ∼ 9.4 cm2g−1, see Hildebrand
1983), and we assume a gas-to-dust ratio of 100. Given the
uncertainties in β and T , we estimate the masses to be accu-
rate within a factor 4. Table 2 gives the derived masses and
beam-averaged column densities. Each sub-peak has a mass
of a few M⊙, and the main submm1 exhibits approximately
10 M⊙ of compact, warm gas and dust emission. The inte-
grated 862 µm continuum flux density of the central region
comprising the four main submm continuum sources amounts
to 1.16 Jy. At an average dust temperature of 100 K, this cor-
responds to a central core mass of 39.9 M⊙. In comparison to
these flux density measurements, Thompson et al. (2006) ob-
served with SCUBA 850µm peak and integrated flux densities
of ∼ 11.5 ± 1.2 Jy/(14′′beam) and ∼19.2 Jy, respectively. The
ratio between peak and integrated JCMT fluxes already indi-
cates non-compact emission even on that scales. Furthermore,
subtracting a typical line contamination of the continuum emis-
sion in hot cores of the order 25% (e.g., NGC6334I, Hunter et
al. in prep.), the total 850µm single-dish continuum flux den-
sity should amount to ∼8.6 Jy. Compared with the integrated
flux density in the SMA data of ∼1.74 Jy, this indicates that
approximately 80% of the single-dish emission is filtered out
by the missing short spacings in the interferometer data. The
dust and gas in the central region have higher temperatures than
the components filtered out on larger spatial scales, and since
the dust and gas mass is inversely proportionally related to the
temperature by MH2 ∝ (e
hν/kT − 1) (e.g., Beuther et al. 2002),
a greater proportion of the mass (> 80%) is filtered out in the
SMA data. However, the SMA image reveals the most compact
hot gas and dust cores at the center of the evolving massive
star-forming region. The shortest baseline of the SMA obser-
vations of ∼16.5 kλ correspond to scales > 12′′ which hence
have to be filtered out entirely. However, even smaller scales are
missing because the uv-spacings corresponding to scales ≥ 5′′
are still relatively poorly sampled and the image presented in
Figure 1 is only sensitive to spatial scales of the order a few
http://cfa-www.harvard.edu/~cqi/mircook.html
4 Beuther et al.: SMA observations of G29.96−0.02
Fig. 1. The hot core UCHii region G29.96−0.02. The grey-scale with contours shows the submm continuum emission with a
spatial resolution of 0.36′′×0.25′′. The contour levels start at the 1σ level of 21 mJy beam−1 and continue at 63, 105 mJy beam−1
(black contours) to 147, 168 mJy beam−1 (white contours). The dashed contours outline the cm continuum emission from the
UCHii region and the thick contours show the NH3 emission (Cesaroni et al. 1994). The contouring is done from 15 to 95% (step
10%) of the peak emission of each image, respectively (S peak(1.2cm) = 109mJy/beam, S peak(NH3) = 15mJy/beam). Triangles,
circles and pentagons show the H2O (Hofner & Churchwell 1996), H2CO (Hoffman et al. 2003) and class ii CH3OH (Minier et al.
2001) maser positions. The star marks the peak of the MIR emission (De Buizer et al. 2002), which is not a point source but has
a similar size as the NH3 emission. The squares mark the infrared sources by Pratap et al. (1999).
arcseconds. The submm peaks detected by the SMA are much
stronger than what would have been expected if the single-dish
flux (∼8.6 Jy) were uniformly distributed over the SCUBA pri-
mary beam of 14′′, even ignoring any spatial filtering and miss-
ing flux effects (This would result in ∼ 4 mJy per synthesized
SMA beam.). This shows that the emission measured on the
small spatial scales sampled by the SMA represents the com-
pact core emission much better than expected. However, it does
not imply that the gas masses measured by the SMA are the
only gas reservoir the embedded protostars have for their on-
going accretion; they may also gain mass from the large-scale
gas envelope that is filtered out by our observations (see also
the competitive accretion scenario, e.g., Bonnell et al. 2004).
The derived beam-averaged H2 column densities are of the or-
der a few times 1024 cm−2, corresponding to visual extinctions
Av of a few 1000 (Av = NH/0.94 × 10
21, Frerking et al. 1982).
3.2. Spectral line emission
Figure 4 presents spectra extracted toward the main submm
submm1 with an angular resolution of 0.64′′×0.47′′ compared
to the submm continuum map (see §2). More than 80 spectral
lines from 18 molecular species, isotopologues or vibrationally
excited species have been identified with a minor fraction of
∼5% of unidentified lines (UL) (Tables 6 & 4). The range of up-
per level excitation temperatures for the many lines varies be-
tween approximately 40 and 750 K (Table 6). Therefore, with
one set of observations we are able to trace various gaseous
temperature components from the relatively colder gas sur-
Beuther et al.: SMA observations of G29.96−0.02 5
Table 2. Submm continuum source parameters
Source R.A. Dec. S peak S int M N
[J2000] [J2000] [ mJy
] [mJy] [M⊙] [10
24cm−3]
submm1 18:46:03.786 -02:39:22.19 173 288 11.5 5.7
submm2 18:46:03.789 -02:39:22.48 168 237 9.5 5.5
submm3 18:46:03.758 -02:39:22.16 138 178 7.1 4.5
submm4 18:46:03.736 -02:39:22.65 151 249 9.9 5.0
submm5 18:46:03.710 -02:39:23.33 68 106 4.2 2.2
submm6 18:46:03.665 -02:39:23.80 84 85 3.4 2.8
The Table shows the peak intensities S peak, the integrated intensities S int, the derived gas masses M as well as the H2 column densities N.
rounding the hot core region to the densest and warmest gas
best observed in some of the vibrationally excited lines.
Table 3. Peak intensities, rms and velocity ranges for images
in Figs. 2 & 3.
Line S peak rms ∆v
mJy/beam mJy/beam km/s
862µm cont., low res. 422 17
CH3OH(73,5 − 62,4) 878 64 [90,104]
13CH3OH(137,7 − 127,6) 752 51 [95,101]
CH3OH(74,3 − 64,3), vt = 1 1419 69 [91,105]
CH3OCH3(74,3 − 63,4) 669 46 [94,104]
C2H5OH(157,9 − 156,10) 586 51 [95,100]
SiO(8 − 7) 391 36 [75,105]
C34S(7 − 6) 592 62 [92,104]
H2CS(101,0 − 91,9) 933 69 [92,100]
34SO(88 − 77) 827 57 [95,103]
SO2(144,14 − 183,15) 544 53 [94,100]
HCOOCH3(275,22 − 265,21) 491 70 [96,100]
CH3CN(198 − 188) 788 71 [94,100]
CH3CH2CN(383,36 − 373,35) 791 56 [94,102]
CH3CHCN(362,34 − 352,32) 655 68 [96,100]
HC3N(37 − 36), v7 = 1 622 55 [94,102]
HC3N(37 − 36), v7 = 2 416 57 [94,100]
HN13C(4 − 3) 1149 76 [94,100]
Figures 2 and 3 now present integrated images of the var-
ious detected species, isotopologues and vibrationally excited
lines. For comparison, Figure 2 also shows the submm con-
tinuum emission reduced with the same degraded spatial res-
olution as the line images. All images show emission in the
vicinity of the hot molecular core and no emission toward the
associated UCHii region. However, the morphology varies sig-
nificantly between many of the observed molecular line maps.
The molecular emission is largely confined to the central region
of the main four submm continuum peaks, and we do not detect
appreciable molecular emission toward the continuum peaks
5 and 6. Reducing the submm continuum data with the same
spatial resolution as the line images, the four submm peaks
are smoothed to a single elongated structure peaking close to
the submm peak submm1 (Fig. 2, top-left panel). The ground
state CH3OH emission is relatively broadly distributed with
two peaks in east-west direction, and one may associate one
with the submm peaks 1 and 2 and the other with the submm
peak submm3, but most other maps show on average one spec-
tral line peak somewhere in the middle of the 4 main submm
continuum peaks, similar to the lower-resolution submm con-
tinuum map.
However, there are also a few species which significantly
deviate from this picture and show a different spatial morphol-
ogy. For example SiO is more extended in north-east south-
west direction likely due to a molecular outflow (§3.3). Also
interesting is the emission from C34S which lacks emission
around the central four submm peaks but is stronger in the in-
terface region between the hot molecular core and the UCHii
region (§4.3). Furthermore, there are a few spectral line maps
– mainly those from likely optically thin lines (HCOOCH3,
HN13C), highly excited lines (CH3CHCN) and vibrationally
excited lines (CH3OH vt = 1, 2, HC3N v7 = 1, 2) – which show
their emission peaks concentrated toward the main submm
peak submm1 (§4.5).
Previous lower-resolution (∼ 10′′) molecular line observa-
tions revealed strong CH3OH emission toward the H2O maser
feature approximately 4′′ south-west of the hot core peak
(Fig. 1, Hofner & Churchwell 1996; Pratap et al. 1999). A lit-
tle bit surprising, we do not detect any CH3OH emission (nor
any other species) toward that south-western position, even
when imaged at low angular resolution using only the com-
pact configuration data (therefore, we do not cover that posi-
tion in Figures 2 and 3). Pratap et al. (1999) discuss mainly
two possibilities to explain this discrepancy: Either their ob-
served specific CH3OH(80 − 71) line is a weak maser and
we do not cover any comparable CH3OH line, or the emis-
sion covered by the lower-resolution data is relatively extended
and filtered out by our observations. As discussed in the pre-
vious section, the shortest baseline of our observations was
∼16 m, implying that we are not sensitive to any scales > 12′′.
Since the CH3OH emission in Pratap et al. (1999) is slightly
resolved by their synthesized beam of 12.6′′ × 9.8′′, it is un-
likely that we would have filtered out all emission. However,
among the many observed CH3OH lines (Table 6), some have
similar excitation temperatures of the order 80 K as the line
observed by Pratap et al. (1999), and we would expect to de-
tect thermal emission from these lines as well. Therefore, our
non-detection of CH3OH emission toward the south-western
H2O maser position supports rather their suggested scenario of
weak CH3OH maser emission in the previously reported obser-
vations (Pratap et al. 1999).
6 Beuther et al.: SMA observations of G29.96−0.02
Fig. 4. Lower and upper sideband spectra extracted toward the
submm1. The spatial resolution of these data is 0.64′′ × 0.47′′.
The main line identifications are shown in both panels.
Table 4. Detected molecular species
Species Isotopologues Vibrational states
CH3OH
13CH3OH CH3OH, vt = 1, 2
CH3OCH3
C2H5OH
HCOOCH3
CH3CN
CH3CH2CN
CH3CHCN
HC3N, v7 = 1, 2
HN13C
a The detection of this CH3OH vt = 2 line is
doubtful since other close vt = 2 lines with
similar excitation temperatures were not detected.
3.3. Molecular outflow emission
The SiO(8-7) spectrum spans a large range of velocities from
∼75 to ∼111 km s−1. Integrating the blue- and red-shifted emis-
sion, one gets the outflow map presented in Fig. 5. The elon-
gated north-west south-east structure is consistent with the pre-
viously proposed outflow by Gibb et al. (2004). The additional
red feature north-east of the central hot core region makes the
interpretation ambiguous: If the north-west south-east outflow
is a relatively highly collimated jet, then the north-eastern red
feature could be attributed to an additional outflow leaving
the core in north-east south-west direction. The blue wing of
that potential second outflow would not detected in our data.
However, since we are filtering out any larger-scale emission,
it is also possible that the red SiO features south-east and north-
east of the main core are part of the same wide-angle outflow
tracing potentially the limb-brightened cavity walls. In this sce-
nario, our observations would miss part of the blue-shifted
wide-angle outflow lobe. With the current data, it is difficult to
clearly distinguish between the two scenarios. However, com-
paring the elongated blue-shifted SiO(8–7) data with the pre-
vious north-west south-eastern outflow observed in H2S by
Gibb et al. (2004), it appears that this is the most likely direc-
tion of the main outflow of the region. Therefore, the multi-
ple outflow scenario appears more likely for the hot core in
G29.96−0.02. The lower resolution SiO(2–1) observation by
Maxia et al. (2001) are also consistent with this scenario. Based
on these data, we cannot conclusively say which of the submm
continuum sources submm1 to submm4 contribute to driving
the outflows.
Fig. 5. SiO(8-7) outflow map. The full and dashed contours
are integrated over the blue- and redshifted SiO emission as
shown in the figure. The contouring starts at ±2σ and continues
in ±1σ steps (thick contours positive, thin contours negative).
The 1σ values for the blue- and red-shifted images are 48 and
46 mJy beam−1, respectively. The markers are the same as in
the previous images, the synthesized beam of 0.68′′ × 0.49′′ is
shown at the bottom right, and the arrows guide the eye for the
potential directions of the two discussed outflows. The offsets
on the axes are relative to the phase center.
4. Discussion
4.1. The formation of a proto-Trapezium system?
The four main submm continuum peaks are located within a
projected area of 7800 × 7800 (AU)2 on the sky. The projected
separation ∆θ between individual sub-sources varies between
1800 AU (peaks 1 and 2) and 5400 AU (peaks 1 and 4, see Table
5). Could the four central submm peaks be the predecessors
of a future Trapezium system? Trapezia are defined as non-
hierarchical multiple systems of three or more stars where the
Beuther et al.: SMA observations of G29.96−0.02 7
largest projected separation between Trapezia members should
not exceed the smallest projected separation by a factor of 3
(Sharpless 1954; Ambartsumian 1955; Abt & Corbally 2000).
This criterion is satisfied by the four submm peaks at the cen-
ter of the G29.96−0.02 hot core. The 14 optically identified
Trapezia discussed by Abt & Corbally (2000) have mean radii
to the furthest outlying member of ∼ 4 × 104 AU, with the
largest radius of ∼ 5.4 × 105 AU (∼2.6 pc), the approximate
dimension of an open cluster. Therefore, the protostellar pro-
jected separations of the tentative proto-Trapezium candidate in
G29.96−0.02 are significantly smaller than in typical optically
visible Trapezia systems. A similar small size for a candidate
Trapezium system has recently been reported for the multiple
system in W3IRS5 (Megeath et al. 2005).
Table 5. Spatial separation
Pair ∆θ ∆x
[′′] [AU]
1-2 0.3 1800
1-3 0.5 3000
1-4 0.9 5400
2-3 0.6 3600
2-4 0.8 4800
3-4 0.6 3600
The numbers in column 1 correspond to the numbers of the submm
peaks.
The small sizes of the proto-Trapezia in G29.96−0.02 and
W3IRS5 may be attributed to their youth. During their upcom-
ing evolution, these young system will expel most of the sur-
rounding gas and dust envelope via the protostellar outflows
and strong uv-radiation. Therefore, the whole gravitational po-
tential of the system will decrease and the kinetic energy may
dominate. Systems with positive total energy will globally ex-
pand and will eventually be observable as a larger-scale optical
Trapezia systems (Ambartsumian 1955).
With the given data it is hard to estimate how massive the
expected Trapezia stars are and will finally be at the end of
their formation processes. The integrated hot core luminosity
is estimated to be ∼ 105 L⊙ (Cesaroni et al. 1994; Olmi et al.
2003), in contrast to the integrated luminosity of the whole re-
gion measured by the large IRAS beam of ∼ 106 L⊙. Producing
105 L⊙ requires either an O7 star or a few stars of comparable
but lower masses. Nevertheless, the numbers imply that this
Trapezium system should form at least one or more massive
stars. Although the gas masses we derived from our dust con-
tinuum data (Table 2) are relatively low, that does not neces-
sarily imply that their mass reservoir is restricted to these gas
masses because it is possible that they may accrete additional
gas from the larger-scale envelope that is filtered out by our
observations. This scenario is predicted by the competitive ac-
cretion model for massive star formation (e.g., Bonnell et al.
2004). The fact that the gas masses we find for the four
strongest submm sources are all similar allows to speculate that
they may form about similar mass stars in the end, however, this
cannot be proven by these data in more detail.
Assuming that the projected size of the potential
proto-Trapezium system in G29.96−0.02 of approximately
7800 (AU)2 resembles a 3-dimensional sphere of radius
∼3900 AU, we can estimate the current protostellar volume
density of the region to approximately 1.4 × 105 protostars
per cubic pc. This number is larger than typical stellar den-
sities in young clusters of the order 104 stars per cubic pc
(Lada & Lada 2003), but it is still below the extremely high
(proto)stellar densities required for protostellar merger models
of the order 106 to 108 stars per cubic pc (Bonnell et al. 1998,
2004; Stahler et al. 2000; Bally & Zinnecker 2005).
Although we have not yet observed the extremely high
(proto)stellar densities predicted by the coalescence scenario,
as soon as we observe massive star-forming regions with a spa-
tial resolution ≤ 4000 AU, we begin to resolve multiplicity and
potential proto-Trapezia (see also the recent observations of
NGC6334I and I(N) by Hunter et al. 2006). Furthermore, this
(proto)stellar density may even be a lower limit, since we ob-
serve only a two-dimensional projection and are additionally
sensitivity limited to masses ≥ 2.1 M⊙ (corresponding to the
3σ flux limit of 63 mJy beam−1 at the assumed temperature of
100 K). Higher spatial resolution has so far always increased
the observed (proto)stellar densities, and it is possible that in
the future we may reach the 106 requirement for merging to
play a role. However, it is also important to get better theoret-
ical predictions of potential merger signatures that observers
could look for.
4.2. Various episodes of massive star formation?
It is interesting to note that the previously identified mid-
infrared source (De Buizer et al. 2002) is offset from the
submm continuum peaks. Although the mid-infrared astrome-
try is usually relatively uncertain, the association of the mid-
infrared peak with class ii CH3OH maser emission with an
absolute positional uncertainty of only 30 mas (Minier et al.
2001) is indicative that the offset may be real. Combining the
facts that we find within a small region of only ∼20000 AU
(∼0.1 pc) at least three different regions of massive star for-
mation – the UCHii region, the mid-infrared source, and the
submm continuum sources – indicates that not all massive stars
within the same evolving cluster are coeval but that sequences
of massive star formation may take place even on such small
spatial scales.
4.3. Carbon mono-sulfide C34S
One of the most striking spectral line maps is from the rare car-
bon mono-sulfide isotopologue C34S(7–6). Its emission peak is
not toward the hot core nor any of the submm continuum peaks,
but largely east of it in the interface region between the submm
continuum peaks and the UCHii region. Hence, one likes to
understand why the C34S emission is that weak toward the hot
core region and that strong at the hot core/UCHii region inter-
face.
CS usually desorbs from dust grains at moderate temper-
atures of a few 10 K, hence it should be observable relatively
8 Beuther et al.: SMA observations of G29.96−0.02
early in the evolution of a growing hot molecular core (e.g.,
Viti et al. 2004). From 100 K upwards H2O is released from
grains, then it forms OH molecules, and the OH can react
with S to SO and SO2 (e.g., Charnley 1997). Therefore, the
initial high CS abundances should decrease with time while
the SO and SO2 are expected to increase with time (e.g.,
Wakelam et al. (2005)). As shown in Figure 2, 34SO peaks to-
ward the hot core where the derived CH3OH temperatures ex-
ceed the H2O evaporation temperature (see §4.4 and Fig. 7b,
potentially validating this theoretical prediction. According to
such chemical models, the hot core G29.96−0.02 should have
a chemical age of at least a few times 104 years.
The strong C34S emission in the hot core/UCHii interface
region may be explained in the same framework. In the molec-
ular evolution scheme outlined above, one would expect low
C34S emission toward the hot core with a maybe symmetrical
increase further-out. In the case of the G29.96−0.02 hot core,
we have the decrease toward the center, but the emission rises
only toward the east, north and west with the strongest increase
in the eastern hot core/UCHii region interface. If one compares
the C34S morphology in Figure 2 with the temperature distri-
bution in Figure 7b, one finds the lowest CH3OH temperatures
right in the vicinity of the C34S emission peaks, adding further
support to the proposed chemical picture.
Extrapolating this scenario to other molecules, it indicates
that species which are destroyed by H2O, e.g., molecular ions
such as HCO+ or N2H
+ (e.g., Bergin et al. 1998), are no good
probes of the inner regions of hot molecular cores.
4.4. Temperature structure
Leurini et al. (2004, 2007) investigated the diagnostic proper-
ties of methanol over a range of physical parameters typical
of high-mass star-forming regions. They found that the ground
state lines of CH3OH are mainly tracers of the spatial density
of the gas, although at submillimeter wavelengths high k tran-
sitions are also sensitive to the kinetic temperature. However,
in hot, dense regions such as hot cores, the effects of infrared
pumping on the level populations due to the thermal heating of
the dust is not negligible, but mimic the effect of collisional ex-
citation. For the ground state line, Leurini et al. (2007) found
that it is virtually impossible to distinguish between IR pump-
ing and pumping by collisions, as both mechanisms equally
populate the vt = 0 levels. On the other hand, the vibrationally
or torsionally excited lines have very high critical densities
(1010–1011 cm−3) and high level energies (T ≥ 200 K). They
are difficult to be populated by collisions and trace the IR field
instead.
To study the physical conditions of the gas around the main
continuum peaks in G29.96–0.02, we analyzed only the emis-
sion coming from the vt = 1 lines, as their optical depth is lower
than for the ground state, and their emission is confined to the
gas around the dust condensations, while the vt = 0 transitions
are more extended and can be affected by problems of missing
flux. We first fitted the methanol emission of the vt = 1 lines
(see Fig. 6) towards the peak position, using the method de-
scribed by Leurini et al. (2004, 2007) that is based on an LVG
analysis and includes radiative pumping (Leurini et al. 2007).
The continuum emission derived in §3.1 was used in the calcu-
lations to solve the equations for the level populations. The two
main dust condensations submm1 and submm2 fall in the beam
of the line data; however, we assumed that the emission is com-
ing from only one component, which is more extended than our
beam, and derived a CH3OH column density averaged over the
beam of 4×1017 cm−2. The corresponding methanol abundance,
relative to H2 is of the order of 10
−7, typical of hot core sources.
Since the emission from the vt = 1 lines is optically thin for
this column density, and also at higher values, we consider this
approach valid. The temperature derived toward the line peak
is 340 K. This corresponds to our best fit model, but from a χ2
analysis we can only infer a low limit of ∼220 K for the temper-
ature of the gas. Since lines are optically thin, the degeneracy
between kinetic temperature and column density is not solved,
and the model delivers good fit to the vt = 1 lines for lower
or higher temperatures by adjusting the methanol column den-
sity. However, the low temperature solutions (Tkin=100–200 K)
need high methanol abundances relative to H2(∼ 10
−6), which
can be hardly found at these temperatures. Moreover, lines are
optically thick for these column densities, and the assumption
of our analysis is not valid anymore.
We also investigated the line ratio between several vt = 1
lines at the column density derived for the main position, to
find the best temperature diagnostic tool among the methanol
lines and derive a temperature map of the region. We found
that the line ratios with the blend of lines at ∼ 337.64 GHz
increase with the temperature of the gas (Fig. 7a). However,
the blending of several transitions together complicates the use
of such diagnostic. In Fig. 7b, we show the map of the line ratio
between the 71,6 → 61,5-E vt = 1 at 337.708 GHz and the blend
between the 71,7 → 61,6-E vt = 1 at 337.642 GHz and 70,7 →
60,6-E vt = 1 at 337.644 GHz. Since line intensities do not
simply add up, we did not correct for the overlapping between
the two transitions. Two other lines, the 74,3 → 64,2-E vt = 1
and the 75,3 → 64,2-E vt = 1, are also very close in frequency.
This is seen in the linewidth of the blending, which is wider
than for the other lines. Therefore, we considered only half of
the channels of the blending at 337.64 GHz in our line ratio
analysis. From the ratio-map in Fig. 7b, submm1, submm2 and
submm3 of Table 1 show high temperatures (T≥ 300 K), while
relatively low temperature gas (T∼ 100 K) is found at R.A.
[J2000]=18h46m03s.818 Dec. [J2000]= −02◦39′22′′.14, close
to a secondary peak of many ground state lines of methanol
(Fig. 2). The temperature then decreases towards submm4. The
increase in the line ratio towards the south-east and north is
probably not true, but due to the poor signal to noise ratio in
these areas. Changes in the column densities along the area may
affect our results.
Fig. 6. Spectrum of the 7ka ,kb → 6ka ,kb−1 vt = 1 methanol band to-
ward the main dust condensation. Overlaid in black is the synthetic
spectrum resulting from the fit.
Beuther et al.: SMA observations of G29.96−0.02 9
Fig. 7. a: Modeled line ratio between the 71,6 → 61,5-E vt = 1 line and
the 71,7 → 61,6-E vt = 1 transitions, as function of the temperature. b:
Map of the line ratio between the same transitions in the inner region
around the peaks. The white stars mark the positions of the dust peaks;
the white dashed contours show the values of the line ratio from ∼ 150
to ∼ 350 K, which correspond to levels from 0.3 to 0.7 in step of 0.1
in the map. The solid black contours show the continuum emission
smoothed to the resolution of the line data (from 0.2 to 0.4 Jy/beam in
step of 0.05). The offsets on the axes are relative to the phase center.
4.5. Tracing rotation toward the massive cores
At the given lower spatial resolution of the spectral line data
compared to the submm continuum, we cannot resolve the four
submm peaks well. However, one of the aims of such multi-line
studies is to identify spectral lines that trace the massive proto-
stars and that are potentially associated with massive disk-like
structures. Such lines may then be used for kinematic gas stud-
ies of rotating gas envelopes, tori or accretion disks. Therefore,
we analyzed the data-cubes searching for velocity structures
indicative of any kind of rotation. In the large majority of spec-
tral lines, this was not successful and we could mostly not
identify coherent velocity structure. While chemical and tem-
perature effects (§4.3 & 4.4) may be responsible for parts of
that, the large column densities derived in §3.1 imply also large
molecular line column densities and hence large optical depths.
Therefore, many of the observed lines are likely optically thick
tracing only outer gas layers of the hot molecular core not pen-
etrating down to the deeply embedded protostars. Furthermore,
many molecules would not only be excited in the central ro-
tating disk-like structures but also in the surrounding envelope
and maybe the outflow. Hence, disentangling the different com-
ponents observationally remains a challenging task.
Fig. 8. Moment 1 maps of HN13C(4–3) (top) and HC3N(37–
36)v7 = 1 (bottom). The markers are the same as in the previous
images, and the synthesized beam of 0.68′′ × 0.49′′ is shown at
the bottom left. The offsets on the axes are relative to the phase
center.
The major exceptions are the molecular lines of the rare
isotopologue of hydrogen isocyanide HN13C(4–3) with a low
excitation temperature of only 42 K, and the vibrationally ex-
cited line of cyanoacetylene HC3N(37–36)v7 = 1 with a higher
excitation temperature of 629 K (Fig. 8). In both cases we find
a velocity gradient across the main submm peak submm1 with
a position angle of ∼ 45◦ from north. This is approximately
perpendicular to the molecular outflow discussed in §3.3 and
by Gibb et al. (2004). Interestingly, Gibb et al. (2004) also find
a similar velocity gradient in their central velocity channels of
H2S. The previously reported NH3 and CH3CN velocity gradi-
ents in approximately east-west direction (Cesaroni et al. 1998;
Olmi et al. 2003) have been observed with slightly lower spa-
tial resolution and are consistent with our data as well.
Our observations as well as previous work in the liter-
ature suggest that the G29.96−0.02 hot core exhibits a ve-
locity gradient in the dense gas in approximately north-east
south-west direction perpendicular to the molecular outflow
observed at larger scales. Based on the HN13C(4–3) map, the
diameter of this structure is ∼ 1.6′′ corresponding to radius
of ∼4800 AU. Since this emission encompasses not only the
submm peak submm1 but also the submm2 and submm3,
it is not genuine protostellar disk as often observed in low-
mass star-forming regions. The velocity structure does not re-
semble Keplerian rotation and may hence be due to some
larger-scale rotating envelope or torus that could transform
into a genuine accretion disks at smaller still unresolved spa-
tial scales (Cesaroni et al. 2007). Additional options to ex-
plain such a velocity gradient may be (a) interaction with the
2nd outflow in north-east–south-western direction, (b) inter-
action with the expanding UCHii region, and (c) global col-
lapse like recently proposed for NGC2264 (Peretto et al. 2006).
While we cannot exclude (a) and (b), option (c) of a globally
collapsing core appears particularly interesting because com-
bining rotation and collapse would result in an inward spi-
raling kinematic structure, potentially similar to the models
originally proposed for rotating low-mass cores (e.g., Ulrich
1976; Terebey et al. 1984). Recent hydrodynamic simulations
by Dobbs et al. (2005) and Krumholz et al. (2006) as well as
analytic studies by Kratter & Matzner (2006) find fragmenta-
tion and star formation within the massive disks forming early
in the collapse process of high-mass cores. This would be con-
sistent with the found three sub-sources (submm1 to submm3)
within the HN13C/HC3N structure. However, on a cautionary
note it needs to be stressed that the collapse/rotation scenario
is far from conclusive, and that the outflow and/or UCHii re-
gion can potentially influence the observed velocity pattern as
well. It remains puzzling that only these two lines exhibit the
discussed signatures whereas all the other spectral lines in our
setup do not.
5. Conclusions and Summary
The new 862 µm submm continuum and spectral line data ob-
tained with the SMA toward G29.96−0.02 at sub-arcsecond
spatial resolution resolve the hot molecular core into sev-
eral sub-sources. At an angular resolution of 0.36′′ × 0.25′′,
corresponding to linear scales of ∼1800 AU, the central core
contains four submm continuum peaks which resemble a
Trapezium-like multiple system at a very early evolutionary
stage. Assuming spherical symmetry for the hot core region,
the protostellar densities are high of the order 1.4 × 105 pro-
tostars per pc3. However, these protostellar densities are still
below the required values between 106 to 108 protostars/pc3 to
make coalescence of protostars a feasible process. Derived H2
column densities of the order a few 1024 cm−2 imply visual ex-
tinctions of a few 1000. The existence of three sites of massive
star formation in different evolutionary stages within a small
region (the UCHii region, the mid-infrared source, and the
submm continuum sources) indicates that sequences of mas-
sive star formation may take place within the same evolving
massive protocluster.
The 4 GHz of observed bandpass reveal a plethora of ap-
proximately 80 spectral lines from 18 molecular species, iso-
topologues or vibrationally excited lines. Only about 5% of the
10 Beuther et al.: SMA observations of G29.96−0.02
spectral lines remain unidentified. Most spectral lines peak to-
ward the hot molecular core, while a few species also show
more extended emission, likely due to molecular outflows and
chemical differentiation. The CH3OH line forest allows us to
investigate the temperature structure in more detail. We find
hot core temperatures≥ 300 K and decreasing temperature gra-
dients to the core edges. The SiO(8-7) observations confirm
a previously reported outflow Gibb et al. (2004) in north-west
south-east direction with a potential identification of a second
outflow emanating approximately in perpendicular direction.
Furthermore, C34S exhibits a peculiar morphology being weak
toward the hot molecular core and strong in its surroundings,
particular in the UCHii/hot core interface region. The C34S de-
ficiency toward the hot molecular core may be explained by
time-dependent chemical desorption from grains, where the
C34S desorbs early, and later-on after H2O desorbs from grains
forming OH, the sulphur reacts with the OH to form SO and
Furthermore, we were interested in identifying the best
molecular line tracers to investigate the kinematics and po-
tential disk-like structures in such dense and young massive
star-forming regions. Most spectral lines do not exhibit any
coherent velocity structure. A likely explanation for this un-
correlation between molecular line peaks and submm contin-
uum peaks is that many spectral lines may be optically thick in
such high-column-density regions, and that additional chem-
ical evolution and temperature effects complicate the picture.
Furthermore, many molecules are excited in various gas com-
ponents (e.g., disk, envelope, outflow), and it is often observa-
tionally difficult to disentangle the different contributions prop-
erly. There are a few exceptions of optically thin and vibra-
tionally excited lines that apparently probe deeper into the core
tracing submm1 better than other transitions. Investigating the
velocity pattern of these spectral lines, we find for some of
them a velocity gradient in the north-east south-west direc-
tion perpendicular to the molecular outflow. Since the spatial
scale of this structure is relatively large (∼4800 AU) compris-
ing three of the central protostellar sources, and since the veloc-
ity structure is not Keplerian, this is not a genuine Keplerian ac-
cretion disk. While these data are consistent with a larger-scale
toroid or envelope that may rotate and/or globally collapse, we
cannot exclude other explanations, such as that the influence
of the outflow(s) and/or expanding UCHII region produces the
observed velocity pattern. In addition to this, these data con-
firm previous findings that the high column densities, the large
optical depths of the spectral lines, the chemical evolution, and
the different spectral line contributions from various gas com-
ponents make it very difficult to identify suitable massive ac-
cretion disk tracers, and hence to study this phenomenon in a
more statistical fashion. (e.g., Beuther et al. 2006)
Acknowledgements. We like to thank Peter Schilke and Sebastian
Wolf for many interesting discussions about related subjects. We
also thank the anonymous referee whose comments helped improv-
ing the paper. H.B. acknowledges financial support by the Emmy-
Noether-Program of the Deutsche Forschungsgemeinschaft (DFG,
grant BE2578).
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12 Beuther et al.: SMA observations of G29.96−0.02
Fig. 2. Compilation of integrated line images (and submm continuum at the same spatial resolution) always shown in grey-scale
with contours and labeled at the bottom of each panel. The dashed contours show negative features due to missing short spacings.
The contouring is done from ±15 to ±95% (step ±10%) of the peak emission of each image, respectively. Peak fluxes S peak, rms
and integrated velocity ranges for each image are given in Table 3. The dotted contours again show the UCHii region and the
stars mark the submm continuum peaks from Figure 1. The offsets on the axes are relative to the phase center.
Beuther et al.: SMA observations of G29.96−0.02 13
Fig. 3. Continued Figure 2.
14 Beuther et al.: SMA observations of G29.96−0.02
Table 6. Line parameters
Freq. Line Eu Freq. Line Eu
GHz K GHz K
337.279 CH3OH(72,5 − 62,4)E(vt=2)
a 727 338.409 CH3OH(70,7 − 60,6)A 65
337.297 CH3OH(71,7 − 61,6)A(vt=1) 390 338.431 CH3OH(76,1 − 66,0)E 254
337.348 CH3CH2CN(383,36 − 373,35) 328 338.442 CH3OH(76,1 − 66,0)A 259
337.397 C34S(7–6) 65 CH3OH(76,2 − 66,1)A
− 259
337.421 CH3OCH3(212,19 − 203,18) 220 338.457 CH3OH(75,2 − 65,1)E 189
337.446 CH3CH2CN(374,33 − 364,32) 322 338.475 CH3OH(75,3 − 65,2)E 201
337.464 CH3OH(76,1 − 60,0)A(vt=1) 533 338.486 CH3OH(75,3 − 65,2)A 203
337.474 UL CH3OH(75,2 − 65,1)A
− 203
337.490 HCOOCH3(278,20 − 268,19)E 267 338.504 CH3OH(74,4 − 64,3)E 153
337.519 CH3OH(75,2 − 65,2)E(vt=1) 482 338.513 CH3OH(74,4 − 64,3)A
− 145
337.546 CH3OH(75,3 − 65,2)A(vt=1) 485 CH3OH(74,3 − 64,2)A 145
CH3OH(75,2 − 65,1)A
−(vt=1) 485 CH3OH(72,6 − 62,5)A
− 103
337.582 34SO(88 − 77) 86 338.530 CH3OH(74,3 − 64,2)E 161
337.605 CH3OH(72,5 − 62,4)E(vt=1) 429 338.541 CH3OH(73,5 − 63,4)A
+ 115
337.611 CH3OH(76,1 − 66,0)E(vt=1) 657 338.543 CH3OH(73,4 − 63,3)A
− 115
CH3OH(73,4 − 63,3)E(vt=1) 388 338.560 CH3OH(73,5 − 63,4)E 128
337.626 CH3OH(72,5 − 62,4)A(vt=1) 364 338.583 CH3OH(73,4 − 63,3)E 113
337.636 CH3OH(72,6 − 62,5)A
−(vt=1) 364 338.612 SO2(201,19 − 192,18) 199
337.642 CH3OH(71,7 − 61,6)E(vt=1) 356 338.615 CH3OH(71,6 − 61,5)E 86
337.644 CH3OH(70,7 − 60,6)E(vt=1) 365 338.640 CH3OH(72,5 − 62,4)A 103
337.646 CH3OH(74,3 − 64,2)E(vt=1) 470 338.722 CH3OH(72,5 − 62,4)E 87
337.648 CH3OH(75,3 − 65,2)E(vt=1) 611 338.723 CH3OH(72,6 − 62,5)E 91
337.655 CH3OH(73,5 − 63,4)A(vt=1) 461 338.760
13CH3OH(137,7 − 127,6)A 206
CH3OH(73,4 − 63,3)A
−(vt=1) 461 338.769 HC3N(37 − 36)v7 = 2 525
337.671 CH3OH(72,6 − 62,5)E(vt=1) 465 338.886 C2H5OH(157,8 − 156,19) 162
337.686 CH3OH(74,3 − 64,2)A(vt=1) 546 339.058 C2H5OH(147,7 − 146,8) 150
CH3OH(74,4 − 64,3)A
−(vt=1) 546 347.232 CH2CHCN(381,38 − 371,37) 329
CH3OH(75,2 − 65,1)E(vt=1) 494 347.331
28SiO(8–7) 75
337.708 CH3OH(71,6 − 61,5)E(vt=1) 489 347.446 UL
337.722 CH3OCH3(74,4 − 63,3)EE 48 347.494 HCOOCH3(275,22 − 265,21)A 247
337.732 CH3OCH3(74,3 − 63,3)EE 48 347.759 CH2CHCN(362,34 − 352,32) 317
337.749 CH3OH(70,7 − 60,6)A(vt=1) 489 347.792 UL
337.778 CH3OCH3(74,4 − 63,4)EE 48 347.842 UL
337.787 CH3OCH3(74,3 − 63,4)AA 48 347.916 C2H5OH(204,17 − 194,16) 251
337.825 HC3N(37 − 36)v7 = 1 629 347.983 UL
337.838 CH3OH(206,14 − 215,16)E 676 348.261 CH3CH2CN(392,37 − 382,36) 344
337.878 CH3OH(71,6 − 61,5)A(vt=2) 748 348.340 HN
13C(4–3) 42
337.969 CH3OH(71,6 − 61,5)A(vt=1) 390 348.345 CH3CH2CN(402,39 − 392,38) 351
338.081 H2CS(101,10 − 91,9) 102 348.532 H2CS(101,9 − 91,8) 105
338.125 CH3OH(70,7 − 60,6)E 78 348.910 HCOOCH3(289,20 − 279,19)E 295
338.143 CH3CH2CN(373,34 − 363,33) 317 348.911 CH3CN(199 − 189) 745
338.306 SO2(144,14 − 183,15) 197 349.025 CH3CN(198 − 188) 624
338.345 CH3OH(71,7 − 61,6)E 71 349.107 CH3OH(141,13 − 140,14) 43
338.405 CH3OH(76,2 − 66,1)E 244
a The detection of this CH3OH vt = 2 line is doubtful since other close vt = 2 lines with similar
excitation temperatures were not detected.
This figure "ch3oh1.jpg" is available in "jpg"
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http://arxiv.org/ps/0704.0518v1
This figure "hc3n_v1_mom1.jpg" is available in "jpg"
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http://arxiv.org/ps/0704.0518v1
http://arxiv.org/ps/0704.0518v1
This figure "hn13c_mom1.jpg" is available in "jpg"
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http://arxiv.org/ps/0704.0518v1
http://arxiv.org/ps/0704.0518v1
This figure "ch3oh2.jpg" is available in "jpg"
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http://arxiv.org/ps/0704.0518v1
This figure "ch3oh3.jpg" is available in "jpg"
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http://arxiv.org/ps/0704.0518v1
	Introduction
	Observations
	Results
	Submillimeter continuum emission
	Spectral line emission
	Molecular outflow emission
	Discussion
	The formation of a proto-Trapezium system?
	Various episodes of massive star formation?
	Carbon mono-sulfide C34S
	Temperature structure
	Tracing rotation toward the massive cores
	Conclusions and Summary
