Muscle w49a data_and_mass_structure_of_the_giant_molecular_cloud

Author version of ApJ, 779, 121
A
Preprint typeset using L TEX style emulateapj v. 5/2/11

MUSCLE W49 : A MULTI-SCALE CONTINUUM AND LINE EXPLORATION OF THE MOST LUMINOUS
STAR FORMATION REGION IN THE MILKY WAY. I. DATA AND THE MASS STRUCTURE OF THE GIANT
MOLECULAR CLOUD.
´
R. Galvan-Madrid1 , H. B. Liu2 , Z.-Y. Zhang3,4 , J. E. Pineda1,5,6 , T.-C. Peng1 , Q. Zhang7 , E. R. Keto7 , P. T. P.
Ho2,7 , L. F. Rodr´
ıguez8,9 , L. Zapata8 , T. Peters10 , C. G. De Pree11

arXiv:1309.4129v2 [astro-ph.GA] 9 Jan 2014

Author version of ApJ, 779, 121

ABSTRACT
The Multi-Scale Continuum and Line Exploration of W49 (MUSCLE W49) is a comprehensive gas
and dust survey of the giant molecular cloud (GMC) of W49A, the most luminous star-formation
region in the Milky Way. The project covers, for the first time, the entire GMC at different scales and
angular resolutions. In this paper we present: 1) an all-configuration SMA mosaic in the 230-GHz
(1.3-mm) band covering the central ∼ 3 × 3 (∼ 10 pc, known as W49N), where most of the embedded
massive stars reside; and 2) PMO 14m telescope observations in the 90-GHz band, covering the entire
GMC with maps of up to ∼ 35 × 35 in size, or ∼ 113 pc. We also make use of archival data from the
VLA, JCMT-SCUBA, IRAM 30m, and the CSO BOLOCAM Galactic Plane Survey. We derive the
basic physical parameters of the GMC at all scales. Our main findings are as follows: 1) The W49
GMC is one of the most massive in the Galaxy, with a total mass Mgas ∼ 1.1 × 106 M within a radius
of 60 pc. Within a radius of 6 pc, the total gas mass is Mgas ∼ 2 × 105 M . At these scales only ∼ 1%
of the material is photoionized. The mass reservoir is sufficient to form several young massive clusters
(YMCs) as massive as a globular cluster. 2) The mass of the GMC is distributed in a hierarchical
network of filaments. At scales < 10 pc, a triple, centrally condensed structure peaks toward the ring
of HC HII regions in W49N. This structure extends to scales from ∼ 10 to 100 pc through filaments
that radially converge toward W49N and its less prominent neighbor W49S. The W49A starburst
most likely formed from global gravitational contraction with localized collapse in a ”hub-filament”
geometry. 3) Currently, feedback from the central YMCs (with a present mass Mcl 5 × 104 M ) is
still not enough to entirely disrupt the GMC, but further stellar mass growth could be enough to allow
radiation pressure to clear the cloud and halt star formation. 4) The resulting stellar content will
probably remain as a gravitationally bound massive star cluster, or a small system of bound clusters.
Subject headings: Hii regions ISM: individual objects (W49A,W49N,W49S,W49SW) stars: formation
stars: massive open clusters and associations: general ISM: clouds galaxies:
starburst galaxies: star clusters
1. INTRODUCTION

Young massive clusters (YMCs) have stellar masses
Mcl > 104 M , sizescales of a few pc, and are younger
than a few crossing times (a few ×100 Myr). They probably represent the young end of the so-called super star
clusters (SSCs) commonly found in starbursting galaxies
1 European Southern Observatory, Karl-Schwarzschild-Str. 2,
85748 Garching, Germany.
2 Academia Sinica Institute of Astronomy and Astrophysics,
P.O. Box 23-141, Taipei 106.
3 Max-Planck Institut f¨ r Radioastronomie, Auf dem H¨ gel
u
u
69, 53121, Bonn, Germany.
4 Purple Mountain Observatory, CAS, 2 West Beijing Road,
Nanjing 210008, China.
5 UK ARC Node, Jodrell Bank Centre for Astrophysics, School
of Physics and Astronomy, University of Manchester, Manchester, M13 9PL, UK.
6 Institute for Astronomy, ETH Zurich, Wolfgang-PauliStrasse 27, CH-8093 Zurich, Switzerland.
7 Harvard-Smithsonian Center for Astrophysics, 60 Garden
Street, Cambridge, MA 02138, USA.
8 Centro de Radioastronom´ y Astrof´
ıa
ısica, UNAM, A.P. 3-72
Xangari, Morelia 58089, Mexico.
9 Astronomy Department, Faculty of Science, King Abdulaziz
University, PO Box 80203, 21589, Jeddah, Saudi Arabia.
10 Institut f¨ r Theoretische Physik, Universit¨t Z¨ rich, Winu
a u
terthurerstrasse 190, CH-8057 Z¨rich, Switzerland.
u
11 Department of Physics and Astronomy, Agnes Scott College, Decatur, GA 30030, USA

through high-resolution imaging with the Hubble Space
Telescope (e.g., Whitmore et al. 1993; Meurer et al. 1995;
Ho 1997). It is also possible that some of them are young
analogues of globular clusters (GCs), provided that they
remain gravitationally bound for timescales comparable
to the Hubble time. Recent reviews are those of Turner
(2009) and Portegies Zwart et al. (2010).
Understanding how YMCs form is not only an interesting problem on itself, but is key to understand the
properties of star forming galaxies elsewhere. In recent
years there have been several studies of deeply embedded YMCs in nearby galaxies that address this problem.
Some studies have targeted the progenitor giant molecular clouds (GMCs) (e.g., Keto et al. 2005; Santangelo
et al. 2009; Wei et al. 2012). Johnson et al. (2001) suggested that the radio continuum sources embedded in
some of these extragalactic GMCs trace the massive star
content, making them analogues to ensembles of ultra
compact (UC) and hypercompact (HC) HII regions embedded in the GMCs that harbor the most luminous star
formation regions in the Milky Way (e.g., Ho & Haschick
1981; Kurtz et al. 1994). However, even in nearby galaxies within 10 Mpc, subarcsecond resolution studies have
physical resolutions not better than ∼ 10 pc.
An alternative approach is to look for embedded YMCs

2

Galvń-Madrid et al.
a

in our Galaxy, which can be well resolved, with the disadvantage that there are only a few of them (e.g., Nguyen
Luong et al. 2011; Liu et al. 2012a; Nguyen Luong et al.
2013; Fukui et al. 2013). Deeply embedded, very luminous (Lbol > 107 L ) star formation regions stand out as
the obvious candidates to be active YMC forming sites.
W49A at Galactic coordinates l = 43.1◦ , b = 0.0◦ is
the most luminous star formation region in the Milky
Way (L ∼ 107.2 L , Sievers et al. 1991), embedded in
one of the most massive giant molecular clouds (GMCs),
Mgas ∼ 106 M (Simon et al. 2001; Miyawaki et al. 2009).
The GMC has an extent of l > 100 pc, but all the prominent star formation resides in the central ∼ 20 pc. This
inner region contains the well known massive star formation regions W49 north (W49N), W49 south (W49S,
∼ 2 southeast of W49N), and W49 southwest (W49SW,
∼ 1.5 southwest of W49N). The most prominent by
far is W49N, hosting the well known ring of HC and
UCHIIs (Welch et al. 1987; De Pree et al. 1997) within
a radius of a few pc. Part of the stellar population in
W49N is already visible in the near-IR and its mass has
been estimated at Mcl 4 × 104 M (Homeier & Alves
2005), whereas the part associated with the most compact HII regions is not even visible in the mid-IR (Smith
et al. 2009a). W49N also hosts the most luminous water
maser in the Galaxy (Gwinn et al. 1992). Those authors determined a direct parallax distance from Earth
of d = 11.4 ± 1.2 kpc. This measurement has been re+0.9
cently improved to 11.1−0.8 (Zhang et al. 2013). We use
the latter value through this paper.
Several possible ideas have been given to explain the
prodigious star formation in W49A. Welch et al. (1987)
proposed that the double-peaked profile seen in most
molecular lines toward the center of W49N is due to
large-scale collapse toward the central ring of HC HII
regions. On the other hand, Miyawaki et al. (1986),
Serabyn et al. (1993), and Buckley & Ward-Thompson
(1996) interpreted it as a cloud-cloud collision. Peng
et al. (2010) showed evidence for expanding shells in the
center of W49N, and proposed that these are the triggering factor for star formation in the whole region.
In this paper we introduce our Multiscale Continuum
and Line Exploration of W49 (MUSCLE W49), a project
aimed at mapping the W49 GMC from its full scale
( 100 pc) down to the scales of individual star-forming
cores ( 0.03 pc). Multiple molecular and hydrogen recombination lines (RLs) are observed, as well as dust
and free-free continuum. When we refer to ”W49” or
the ”W49 GMC” we mean the full W49 giant molecular
cloud, when we use ”W49A” we refer to the central part
of the GMC that is actively forming YMCs. When necessary, we refer individually to the common labeling of the
subcomponents of W49A: W49N, W49S, and W49SW.
This paper presents the bulk of the data set and derives
the mass structure of the cloud at all scales. Upcoming
papers will deal with further scientific analysis: a multiscale analysis of the dynamics, a comparison of hot cores
and hypercompact HII regions using the subarcsecond
resolution mosaics of W49N, and a quantitative comparison with extragalactic star formation. Section 2 and Appendix A describe the observations. Section 3 presents
the data. In Section 4 we derive the mass structure of
the GMC, and in Section 5 we discuss on the GMC and

cluster disruption. We list our conclusions in Section 6.
Appendices B and C show the rest of the line maps from
the PMO and SMA observations, respectively. Appendix
D describes our method to obtain the gas surface density
maps using CO and its isotopologues, and to derive gas
masses from the dust continuum.
2. OBSERVATIONS
2.1. PMO 14m

Mapping observations were made from March 2011 to
July 2011 with the 14-m telescope of Purple Mountain
Observatory12 (hereafter PMO 14m), located at Delingha, Qinghai, China. The CO and its isotopologues
were observed in the 1–0 transition, as well as several
other molecular lines. The observations were performed
with a nine pixel SIS receiver, which was configured with
dual sideband (2SB) mixers for each pixel (Shan et al.
2010). For each sideband and each pixel, fast fourier
transform (FFT) spectrometers produce a bandwidth of
1 GHz and 16384 channels, resulting in a velocity resolution of ∼0.16 km s−1 and a velocity coverage of ∼ 2700
km s−1 at 110 GHz.
The mapping observations were carried out in on-thefly (OTF, e.g., Mangum et al. 2007) scan mode. A position of 120 to the north of the centre of W49A was
taken as sky reference. All maps are done with uniform
Nyquist sampling, except for ∼2 around the edges. The
rms pointing uncertainty is estimated to be better than
5 . Typical system temperatures during the runs were
∼ 140 K at 110 GHz and ∼ 270 K at 115 GHz.
All the PMO spectral line data were reduced with
the CLASS/GILDAS13 package developed by IRAM. We
classify the spectral quality of each spectrum by the baseline flatness and system temperature levels. About 5% of
the spectra are discarded due to poor baselines. Linear
baselines are subtracted for each individual spectrum.
All spectra are then co-added and re-grided.
We converted the antenna temperature (TA ) to the
main beam brightness temperature (Tmb ) scale using
Tmb = TA (ηf /ηmb ), where the ratio of the main beam
efficiency ηmb to the forward hemisphere efficiency ηf is
0.46.
For each molecular line, we combined all the calibrated
spectra and re-gridded them to construct datacubes with
weightings proportional to 1/σ 2 , where σ is the rms noise.
This routine convolves the gridded data with a Gaussian
kernel of ∼1/3 the telescope beamsize, yielding a final
angular resolution slightly coarser than the original beam
size. The HPBW at the CO frequency is 58.0 (3.4 pc).
2.2. SMA

We observed the central cluster of W49A (W49N) with
the Submillimeter Array14 (Ho et al. 2004) in the 1.3mm (230-GHz) band using the four different array configurations: subcompact, compact, extended, and very
extended, one track for each. This multi-configuration
approach images both small and large scales. The data
12

http://www.dlh.pmo.cas.cn/
http://www.iram.fr/IRAMFR/GILDAS
14 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.
13

MUSCLE W49. Data and Mass Structure.
set covers baseline lengths between 8 kλ (corresponding
to scales of 31 ) and 480 kλ (0.5 ). Figure A1 (appendix
A) shows the (u, v) coverage of the combined data set.
The upgraded SMA correlator capabilities were used,
covering 2 × 2 GHz in each of the two sidebands. The
lower sideband (LSB) covered the sky frequency ranges
from 218.29 to 220.27 GHz in spectral windows 48 to 25,
and from 220.29 to 222.27 GHz in spectral windows 24
to 1. The upper sideband (USB) covered from 230.29
to 232.27 GHz in spectral windows 1 to 24, and from
232.29 to 234.27 GHz in spectral windows 25 to 48. The
spectral resolution of the raw data set is ≈ 1.1 km s−1 .
The visibilities of each track were separately calibrated
using the SMA’s data calibration program, MIR15 . The
phase, bandpass, and flux calibrators for each track are
listed in Table 1. Imaging and basic analysis were done
in CASA versions 3.3 and 3.4.
The continuum emission was subtracted in the (u, v)
domain by fitting a baseline across the passband using
line-free channels as input. The continuum and line data
of each configuration were then combined into a single
data set and imaged. The concatenated continuum data
was self-calibrated, then imaged with different weights,
checking that the flux recovered at different resolutions
was the same within a few percent. The continuum map
presented in this paper is the result of the combination of
the SMA mosaic with the CSO BOLOCAM GPS archival
map (Ginsburg et al. 2013). The combination procedure is based on converting the single-dish map to the
Fourier space, and then jointly inverting and cleaning the
combined single-dish and interferometer visibilities. The
13
CO and C18 O SMA mosaics were also combined with
single dish maps to fill the short spacings. We used the
IRAM 30m maps presented in (Peng et al. 2010, 2013).
The combining procedure is described in more detail by
Liu et al. (2013).
The 11-pointing SMA mosaic is Nyquist sampled and
covers all of the HC HII regions detected in the cm by De
Pree et al. (1997, 2004). All the presented SMA mosaics
have been corrected for primary-beam attenuation. Figure A1 (appendix A) shows the primary-beam response
of the mosaic used to correct for attenuation in the final
images.
Archival maps from JCMT-SCUBA16 at 678 GHz
(project code M97Bu89) and the Very Large Array17 at
8.5 GHz (project code AD324) are also used in this paper.
We refer the reader to Di Francesco et al. (2008) and De
Pree et al. (1997) for the respective observational details.
Further analysis of the above mentioned data sets
was performed in CASA18 , MIRIAD (Sault et al. 1995),
GILDAS19 , DS920 , Karma (Gooch 1996), IDL, and
Python.

3

3. PRESENTATION OF THE DATA

We present the data zooming in from large to small
scales.
3.1. The W49 Giant Molecular Cloud
3.1.1. CO 1–0 and isotopologues

The largest-scale maps are those obtained with the
PMO 14m telescope in CO J = 1 − 0 and its isotopologues 13 CO and C18 O. These maps cover the entire W49
GMC with up to 37 (119 pc) per dimension. Figure 1
shows the velocity-integrated (moment 0) CO emission
maps. This emission covers the LSR velocity range from
−20 km s−1 to 30 km s−1 . Outside this range (up to
VLSR = 78 km s−1 ) there are more CO spikes. Their
line profiles are narrow and their emission cover the entire mapped area without matching the features from the
W49 GMC. We infer that they are clouds in the line of
sight not associated with W49.
The CO maps show the well-known central gas clump
in the inner 3 , known as W49 North (W49N). It is in this
region where the central embedded cluster with dozens
of massive stars as traced by radio-continuum hypercompact (HC) and ultracompact (UC) HII regions reside (De
Pree et al. 1997, 2004). The central W49N clump appears
connected to the W49 GMC via a series of filamentary extensions (see Fig. 1). These filaments are hinted in previous observations where they appear as protuberances
out of W49N toward the southwest, the east-southeast,
and the north (see, e.g., fig. 2 of Nagy et al. 2012, or
fig. 6 of Matthews et al. 2009). The southeast filament
connects W49N to W49 South (W49S) located 2 apart,
and which hosts a cometary HII region (Dickel & Goss
1990).
The spectrum from the W49 GMC has towards its position centroid two prominent velocity features. Figure 2
shows the CO spectra toward the moment 0 peak. The
optically thick 12 CO is too complex to be fitted with a
small number of Gaussians. The isotopologues are well
fit by a sum of three Gaussian components. The kinematics of the full GMC and its relation to smaller scales
will be the topic of a following paper.
3.1.2. Other molecules and hydrogen

Figures B1 and B2 in Appendix B show the moment 0
(integrated from −12 km s−1 to 24 km s−1 ) of the rest of
the molecules clearly detected in the PMO 14m observations, as well as of the hydrogen recombination line (RL)
41α. Figures 3 and 4 show their respective spectra at the
integrated-emission peak. All the transitions peak at the
same position within a few arcseconds, a fraction of the
beam. Table 2 lists all the lines detected with PMO and
their velocity-integrated fluxes spatially integrated over
the projected area of the GMC. The intensity conversion
15 https://www.cfa.harvard.edu/ cqi/mircook.html
factor at the rest frequency of the CO 1–0 is 0.027 K per
16 The James Clerk Maxwell Telescope is operated by the Joint
Jy beam−1 .
Astronomy Centre on behalf of the Science and Technology FacilTable 3 gives the parameters obtained from Gaussian
ities Council of the United Kingdom, the Netherlands Organisafitting to the cloud components for the different line traction for Scientific Research, and the National Research Council of
Canada.
ers. The bright (peak amplitude A > 1 K) lines of HCN
17 The National Radio Astronomy Observatory is a facility of
1–0 and CS 2–1 follow the CO isotopologues in havthe National Science Foundation operated under cooperative agreeing at least two clear velocity features peaked at ∼ 3
ment by Associated Universities, Inc.
18 http://casa.nrao.edu/docs/UserMan/UserMan.html
km s−1 and ∼ 12 km s−1 (the first peak of the CO19 http://www.iram.fr/IRAMFR/GILDAS/
isotopologues is at ∼ 4 km s−1 ). The bright HCO+ 1–0
20 https://hea-www.harvard.edu/RD/ds9/site/Documentation.html
appears to have its two components slightly blueshifted

4

a

to 2 and 11 km s−1 respectively. The faint isotopologues
H13 CN 1–0 and H13 CO+ 1–0, as well as the N2 H+ 1–
0, also have two components but their parameters have
larger uncertainties. The SiO 2–1, SO 2(2)–1(1), and
SO2 8(1,7)–8(0,8) were best fit by a single Gaussian. The
signal-to-noise in the latter two is too low to distinguish
two Gaussians. The ionized hydrogen as traced by the
H41α RL appears single peaked at 12 km s−1 . As expected, this line is far broader than the molecular lines
(FWHM≈ 30.2 km s−1 ).
Although we use several Gaussians to quantify the line
features, we do not interpret them as evidence for separate clouds. From the CO spectra it is clear that most
of the double-peaked feature is caused by self absorption
due to the very high column densities in the line of sight
to the cloud center and around ∼ 7 km s−1 . This is confirmed by our calculations to create the surface density
maps of Section 4.1. However, there is also a velocity
gradient across W49N that contributes to create the line
asymmetry.
3.2. The Scale of the Central Clusters: W49N and

W49S
At the center of the W49 GMC lies the well known
massive star formation region W49A, the most luminous
in the Milky Way. W49A has two main clusters separated by ∼ 2.5 (∼ 8 pc): W49N (the most prominent
with dozens of deeply embedded (possibly still forming) O-stars as traced by UC and HC HII regions De
Pree et al. 1997), and W49S (to the southeast). These
scales are mapped at subarcsecond angular resolution,
but with sensitivity to scales as large as 20 with the
all-configuration SMA mosaic covering the inner ∼ 2.5 .
The continuum mosaic, as well as the 13 CO and C18 O
2–1 mosaics are combined with single-dish maps to recover the missing extended emission. Archival single-dish
(sub)mm maps, as well as Very Large Array (VLA) cm
maps that cover both W49N and W49S are also used.
3.2.1. The Continuum Emission

Figure 5 top shows the archival CSO BOLOCAM
Galactic Plane Survey image (Aguirre et al. 2011; Ginsburg et al. 2013) at 1.1 mm of the central ∼ 25 × 25 pc of
the W49 GMC. It covers W49N, W49S and filamentary
extensions in between and toward them. The total flux
in the BOLOCAM image is 101.0 Jy.
It is known that HC and UC HII regions can have
α
rising spectral indices (αff , where the flux Sν ∝ νff ) in
their free-free emission from cm- all the way up to mm
wavelengths, with values αff ∼ 0.5. This is due to their
high densities and non-zero optical depths (Keto et al.
2008; Galvń-Madrid et al. 2009; Peters et al. 2010b).
a
Therefore, to obtain dust masses when young HII regions are present it is necessary to subtract the free-free
contribution to the mm flux. We use the 3.5-cm (8.5
GHz) and 7-mm (45.5 GHz) VLA data presented by De
Pree et al. (1997, 2004) to separate the total ionized-gas
emission from the total dust emission. The 3.5-cm map
covers both W49N and W49S and matches the BOLOCAM map. The 7-mm map covers only W49N where
most of the HC sources reside. The total free-free flux
at 8.5 GHz is 29.7 Jy. The compact sources have an effective αff = 0.2, and contribute with 4.4 Jy of flux at 7

mm and 3.2 Jy at 3.5 cm, only 11% of the total flux at
the larger wavelength. We take an effective αff = 0 for
the free-free flux from 3.5 cm to 1 mm. Therefore the
contributions at 1 mm are ∼ 29% free-free, and ∼ 71%
dust (71.3 Jy in the BOLOCAM map). Feasible values of
the effective αff of the compact+extended emission are
limited between -0.1 (optically thin emission) and 0.1 (a
few compact sources may be outside W49N), then the
limits to the dust contribution at 1 mm are 58% to 79%.
We now proceed to estimate the total gas mass in the
central clusters (W49N+W49S) from the dust emission.
For this we use equation D6 of Appendix D and a gasto-dust mass ratio of 100. Studies of the SED of W49N
from the mid infrared to the millimeter have found dust
temperatures Td from 20 K to 50 K (Ward-Thompson &
Robson 1990; Sievers et al. 1991). We assume an effective Td = 30 K. For Td in the extremes of the above mentionted range, the derived masses change by about +70%
and −45%. However, the main uncertainty in deriving
masses from dust emission comes from the assumed dust
absorption coefficient κν . For high-mass star formation
regions, values in the range from 0.1 to 0.5 cm2 g−1 (per
unit mass of dust) at 1.1-mm are appropriate (Ossenkopf
& Henning 1994; Galvń-Madrid et al. 2010; Beuther
a
et al. 2011). The total gas mass that we obtain from 1mm dust emission is Mgas,d ∼ 1.5×105 to 7.6×105 M for
Td = 30 K. The lower end of this estimation (higher dust
opacity) agrees with the CO measurements presented in
Section 4. The estimations are also within the range of
values obtained by Nagy et al. (2012) and Peng et al.
(2013).
Figure 6 shows an archival JCMT-SCUBA map at 678
GHz (0.4 mm) of the central clusters W49N+W49S. The
angular resolution of this map is 8.1 , intermediate between the combined SMA+BOLOCAM and the BOLOCAM images. Because dust is so much brighter at higher
frequencies, this map is sensitive enough to image the
filamentary structures joining W49N to W49S, and extending to the southwest (sometimes called W49SW),
southeast, and north. The magnetic field in W49 has
been detected by Curran & Chrysostomou (2007) using
polarized submillimeter dust continuum emission. The
plane-of-the sky magnetic field is relatively weak (less
than 0.1 mG) and well aligned along the filament joining
W49N to W49S. In the rest of the paper we show that
the filamentary structures are also seen in CO integrated
intensity and column density maps, and that continue all
the way up to the scale of the full GMC (> 100 pc).
The dust mass can be independently derived from the
0.4 mm emission. Note that at 678 GHz and for Td = 30
K we are completely out of the Rayleigh-Jeans regime
(hν/kT ∼ 1), so we have to use the full Planck function as in equation D6. At this wavelength there is no
significant contribution from free-free, and the values of
the dust absorption coefficient κν at shorter wavelenghts
are less sensitive to the properties of dust (e.g., see Fig.
5 in Ossenkopf & Henning 1994). From Ossenkopf &
Henning (1994), we consider the range of κ0.4mm values for coagulated dust grains without thick ice mantles, 6.5 to 9.8 cm2 g−1 . Then the flux in the 0.4-mm
map S0.4mm = 1.96 × 104 Jy corresponds to a gas mass
M = 4.2 × 105 to 6.3 × 105 M .
We now describe the features of the high angular reso-

lution SMA+BOLOCAM mosaic of W49N (figure 5 bottom, covering the region marked by contours in Figure
5 top). The combining procedure is based on creating
a visibility data set from the single-dish map that then
is combined with the interferometer visibilities, fouriertransformed back to the image plane, and jointly cleaned,
as described in Liu et al. (2012b, 2013). The combined SMA+BOLOCAM map has an angular resolution
of ∼ 2.3 . The subarcsecond resolution dust maps from
the SMA mosaic only, which do not recover the extended
emission but are sharper will be discussed along with the
subarcsecond resolution SMA line mosaics in a future article. In this paper we focus on the mass structure of the
W49N cluster as a whole.
To quantify the structure in the SMA+BOLOCAM
continuum map, we prefer to measure the total mass in
the maps above a series of intensity (mass surface density) thresholds, rather than trying to extract sources
using source-finding algorithms. The reason is that
these algorithms tend to artificially fragment the emission (Pineda et al. 2009). If one acknowledges the fact
that the ISM is hierarchical at all scales (except possibly
at < 0.1 pc where structures will end up forming a single star/binary system), it follows that trying to count
sources is not very meaningful. We created masked maps
with intensities > 5 mJy beam−1 ×5, 10, 20, 40, 80, 160,
and 320 (the peak in the map is 3.302 Jy beam−1 ) and
measured their total flux. The positive contour in Figure
5 bottom is at 25 mJy beam−1 . Table ?? lists the results.
The quoted values for surface densities and dust-derived
masses correspond to κ1mm = 0.25 cm2 g−1 , and the errors to the results from varying κ1mm from 0.1 to 0.5 cm2
g−1 . As mentioned before, variations in Td introduce an
extra uncertainty (see Appendix D for details). We have
subtracted a free-free contribution of 58 %21 .
3.2.2. The Line Emission

To identify the spectral features we first imaged the full
bandwidth covered by the SMA using the subcompact
array data only. Figure 7 shows these spectra spatiallyaveraged in a contour matching the 5σ 1.3-mm continuum shown in Figure 5, bottom. The molecules/atoms
producing the line emission are labeled. The details
on the species, transition, rest frequency, and upperlevel energy are listed in Table 5. All the 57 lines that
were clearly identified (including some blended features)
above an intensity threshold of 50 mJy beam−1 are listed.
Many of the detected lines are typical ’hot-core’ tracers
characteristic of luminous star formation regions (e.g.,
Kurtz et al. 2000; Cesaroni et al. 2010; Beltrń et al.
a
2011), although W49A appears to be less line-rich than
its cousin Sgr B2 near the Galactic Center (e.g., Nummelin et al. 1998; Belloche et al. 2013). We have unambigously identified 40 individual lines, plus 14 lines that
are clearly identified but their profiles are blended (e.g.,
the CH3 CN J = 12 − 11 K = 0 and K = 1 lines whose
centers are only separated by only 5.7 km s−1 ). Three
more identified lines possibly have significant contamination (see Table 5).
21 As derived before, the total free-free contribution in the entire
VLA map is ∼ 21%. Within the area of the SMA+BOLOCAM
map (W49N), an estimation of the free-free contribution under the
same assumptions gives ∼ 58%.

5

W49A appears to be rich in SO2 and its isotopologues.
Younger massive star formation regions devoid of ionization or with only faint HII regions tend to lack emission in these lines (e.g., Galvń-Madrid et al. 2010; Liu
a
et al. 2012a), whereas regions in which O-stars have already formed and with HC and UC HII regions tend to
be brighter in the SO2 lines in this band (e.g., Galvńa
Madrid et al. 2009; Baobab Liu et al. 2010).
After line identification, individual line cubes with the
all-configuration SMA data were created. The synthesized beam of these maps varies smoothly from 1.02 ×
0.69 , PA= 76.8◦ at the lowest frequency (36.4 K per
Jy beam−1 ) to 0.98 × 0.64 , PA= 74.1◦ at the highest (31.5 K per Jy beam−1 ). The resolution of the final
maps was chosen such that nearly all of the flux of the
subcompact-array maps is recovered but at the highestpossible angular resolution. The CO, however, is affected
severely by missing flux in interferometric observations.
We have therefore combined the multiconfiguration SMA
maps with IRAM 30m maps for the 13 CO 2–1 and C18 O
2–1. Figure 8 shows the velocity-integrated (moment 0)
of these two CO isotopologues. As with the larger-scale
PMO data, the emission has been integrated in the LSR
velocity range from −20 km s−1 to 30 km s−1 . The angular resolution of the single-dish+interferometer maps
is ≈ 2 .
A striking feature of the SMA+IRAM30m maps is that
the radially-converging, triple filametary structure seen
at the full GMC scales (from ∼ 10 to 100 pc) as traced
by the PMO maps (section 3.1) is preserved all the way
down to the inner few pc. In Section 4 we discuss in more
detail the morphological matching between these structures and its physical interpretation. Their dynamics will
be discussed on a following paper.
Appendix C (on-line only) shows the velocityintegrated (moment 0) line emission of all the SMAdetected lines listed in Table 5. A single map is shown
for neighbouring lines whose profiles are blended. The
last column of Table 5 lists the velocity-integrated line
fluxes. For the science discussion of this paper (sections 4
and 5) we use the combined SMA+IRAM 30m CO maps
together with the dust maps and the larger-scale PMO
CO maps to make a study of the mass of the W49 GMC
as a function of scale.
4. THE MASS STRUCTURE OF THE W49 GMC
4.1. Mass surface density maps

We calculate maps of the mass of the GMC from scales
of ∼ 1 pc to > 100 pc using the PMO CO maps on
the larger scales and the SMA+IRAM 30m CO maps
on the smaller scales. The methodology is described in
Appendix D, and it is based on solving for the optical
depth in every voxel of data from the CO-isotopologue
ratios. Uncertainty in this mass determination comes
from uncertainties in the isotopologue abundances (typically ±50%), in the CO to H2 conversion factor, and
in the assumed excitation temperature Tex (typically
±25%). All added up, masses are uncertain within a
factor of 2 each way (see Appendix D). Figure 9 shows
the resulting maps of molecular-gas surface density Σ in
units of M pc−2 . We derive a total mass for the GMC
Mtot = 1.14 × 106 M . In the area covered by the 0.4mm dust map (Section 3.2.1), the CO mass map gives a

6

a

mass MW49N+S+SW = 4.1 × 105 M , in agreement with
the lower values in the range of masses as derived from
0.4-mm dust emission (M = 4.2 × 105 to 6.3 × 105 M ).
On the full-GMC map (top frame) the central cluster
W49N is the most prominent feature in the center. Radially out of it, the known triple filamentary structure of
scale ∼ 30 pc is seen connecting W49N to W49S (to the
southeast), to W49SW (to the southwest), and extending north of W49N. W49S appears to be connected to a
clump ∼ 4 (∼ 13 pc) to the northeast22 . W49SW appears to end in a high density peak. Another high column
density clump is seen ∼ 5.5 (∼ 18 pc) southwest of the
main W49N peak, and is joined by a north-south extension to the brighter filament joining W49N to W49SW.
A chain of filaments extend north of W49N and all of
them appear curved with a similar orientation. The full
GMC mass distribution is clumpy, which indicates fragmentation also along the filaments in the periphery. All
the bright UC and HC HII regions reside in the central
main clusters, and W49N is by far the richest in ionizedgas structures (De Pree et al. 1997, 2004). A possible
future line of research would be to study star formation
in detail in the rest of the high column-density peaks
as it has been done with W49N, looking for fainter HII
regions and molecular-gas cores.
The bottom panel of Figure 9 shows the high-angular
resolution column density map obtained from the combination of the CO SMA mosaics with the IRAM 30m
observations presented in Peng et al. (2013). These map
cover the area marked by a black square on the fullGMC map. The global column density peak is right in
the middle of the ring of HC HIIs (see Figure 8 for an
overplot), at the position of the HCHII labeled as B by
De Pree et al. (1997). The highest column density features (Σ > 104 M pc−2 ) are indeed the highest volume
density regions: in Figure C1 of Appendix C it is seen
that many of the molecular lines that trace the densest
gas (n > 105 cm−3 ) peak at the same position as the
high-resolution CO column density maps.
It is striking to notice that the larger scale (10 to 100
pc) structure of the cloud is preserved all the way to
the inner < 10 pc. This triple structure consists of one
structure of filaments oriented toward the east-southeast
(which join W49N to W49S), one structure of filaments
toward the southwest, and filaments extending to the
north. Each of the filamentary structures is seen to be
composed of a tree of filaments in itself.
4.2. A Hub-filament network
The maps reveal that the mass structure of the W49
GMC is organized in a hierarchical network of filaments
that appear to converge in the central, densest region
harboring the W49N and W49S clusters. This network
extends all the way from the inner pc where the embedded O-type stars in W49N reside (as traced by the dozens
of UC and HC HII regions, see Figure 8, or Figure 1 in De
Pree et al. 1997) to the > 100 pc scales of the full GMC
(see Figure 1). It is also preserved at intermediate scales
(∼ 10 pc), where the filaments connect W49N to W49S,
and both of them to the surrounding GMC (figures 5 and
6).
22 This clump is also seen in archival single-dish dust images and
in the individual CO maps (figure 1).

This type of hierarchical, centrally-condensed network
of filaments appears to be common in clouds harboring the most luminous (L > 105 L ) star formation
regions in the Galaxy. For example, the global structure of the W49 GMC appears to be a scaled up version of the GMC in the G10.6–0.4 massive star formation region (L ∼ 9.2 × 105 L , see Figure 3 of Liu et al.
2012b). In this scenario, the densest, central regions of
the GMC collapse first because they have shorter freefall time (tff ∝ ρ−1/2 ). Furthermore, departures from
spherical to sheet or filamentary geometries increase the
free-fall time with the aspect ratio of the configuration
(Toal´ et al. 2012; Pon et al. 2012). Although never
a
completely spherical, the central W49N clump is closer
to sphericity than the larger-scale filaments23 . Therefore,
it is natural to expect that the central clump (scales < 10
pc) with an average number density n ∼ 104 to 105 cm−3
can collapse and form massive stars in a time-scale of the
order of a few times the spherical-collapse free-fall time
t ∼ 105 yr, whereas the more filamentary, larger-scale
GMC, even with an initial density close to that of the
central clump, will collapse in a timescale > 1 Myr. Figure 1 of Liu et al. (2012b) shows a schematic diagram
that accounts for the aspects above discussed.
Filament networks are often also present in low- and
intermediate-mass star formation regions (Gutermuth
et al. 2008; Myers 2009; Molinari et al. 2010; Pineda et al.
2011; Kirk et al. 2013; Takahashi et al. 2013). However, it
is unclear why some of them look like a filament-only network and some other also have prominent hub-like condensations. This could be due to evolution/mass differences, or due to different observational techniques24 . We
hypothesize that centrally condensed networks are the
main morphology in the most massive clusters born out
of the most massive clouds like W49 or Carina (Preibisch
et al. 2012; Roccatagliata et al. 2013), but more systematic studies are needed. In W49, for example, many
previous high-resolution studies focused only on the central few pc of the main clusters (e.g., Wilner et al. 2001;
Peng et al. 2010), whereas studies that captured the entire GMC lacked the sensitivity to recover the external
filaments (e.g., Simon et al. 2001; Matthews et al. 2009).
Hierarchical filament networks surrounding a central
cluster (or small system of clusters) are also naturally
expected from simulations of massive clouds that form
massive star clusters (e.g., Bonnell et al. 2001; Smith
et al. 2009b; V´zquez-Semadeni et al. 2009; Dale et al.
a
2012). In these simulations, the primorial GMC has density inhomogeneities (presumably caused by the process
of GMC formation itself) that give rise to filamentary
structure while the full cloud contracts. These filaments
can then converge25 into the hubs where cluster formation occurs (e.g., Galvń-Madrid et al. 2010; Duartea
Cabral et al. 2011; Inoue & Fukui 2013), preferably at
locations deep within the global gravitational potential.
23 i.e., the filament-interfilament column density contratst is
larger in the cloud periphery.
24 The dust continuum, in contrast to line observations, picks
everything on the line of sight.
25 Although less used in the literature, we prefer the word ”convergence” to ”collision”, since the latter implies clearly distinct
entities, and does not apply to, for example, a case where different
filamentary arms that originated from a single GMC converge to a
point.

When feedback from the formed stars is included, it is
usually found that it has an important effect in regulating the star formation efficiency (e.g., Peters et al. 2010a,
2011; Krumholz et al. 2011; Dale et al. 2012; Col´ et al.
ın
2013).
4.3. Other interpretations
Several interpretations have been put forward to explain the observations of W49N. Welch et al. (1987)
proposed a global collapse scenario based on interferometric observations of RLs of the ring of HC HIIs and
molecular-line absorption toward their line of sight. Keto
et al. (1991), based on modeling the line profiles resulting from a hydrodynamical simulation, showed that the
observations of Welch et al. (1987) could be explained
by a global-local collapse scenario where fragmentation
from the larger-scale contracting cloud produces denser
molecular cores that themselves collapse into individual
HII regions. Serabyn et al. (1993) proposed that the
multiply-peaked single-dish line profiles are due to different molecular clumps whose collision triggered the vigorous star formation activity. The morphological evidence
presented in this paper shows that the most likely situation is similar to a scaled-up version of the Keto et al.
(1991) model. The filamentary ∼ 100 pc scale GMC
has fragmented and the denser W49N clump (extending
further toward the southeast, southwest, and north) is
the converging region of a hierarchical (triple) network
of filaments that converge toward the ring of HC HIIs.
If the different filaments are viewed as separate entities
then they could be interpreted as colliding clouds. However, the observations presented in this paper show that
they are part of a larger, common structure. Kinematical evidence supporting this interpretation will be shown
in the second paper of this project. Further evidence
comes from looking at the available single dish line profiles: the larger the optical depth, the more prominent
the ∼ 7 km s−1 dip is, indicating that it is produced
by self-absorption due to the large column of gas toward
the center of W49N. Indeed, we obtain our larger optical depths from the line ratios at these positions and
velocities.
However, the ring of HC HIIs is not the entire story.
There is already a NIR massive star cluster concentrated
within 1 pc and located ∼ 3 pc east of the Welch ring
(cluster 1 of Alves & Homeier 2003). This cluster likely
is the ionizing source of the extended halo of free-free
emission next to the HC HIIs (see Figure 8), whose edge
matches the shell reported by Peng et al. (2010) from
4.5 and 8.0 µm observations. These two YMCs (which
may be parts at different evolutionary stages of the same
YMC), together with the rest of the young massive stars
in W49N, W49S, and W49SE, formed from the same
GMC.
4.4. Comparison to other massive clouds in the Galaxy

We now proceed to use the column density maps described in Section 4.1 to compute the mass of the GMC
as a function of distance from the central cluster W49N,
and to compare with other massive GMCs in the Galaxy
that are believed to be progenitors of young massive clusters (YMCs). Figure 10 shows the result. The match between the calculation using the CO 1–0 ratios on larger

7

scales overlaps remarkably well with the result on smaller
scales using CO 2–1 ratios. The mass in a region around
W49N with a diameter D = 8.0 pc M ∼ 1.2 × 105 M ,
equivalent to a a mean density under the assumption of
spherical symmetry of ∼ 9 × 103 cm−3 . Likely, the mean
density is higher because of the filamentary geometry and
substructure. Indeed, summing the CO mass in an aperture matching the dust filaments seen at 0.4 mm gives
M ∼ 4 × 105 M . The full GMC on scales of D ∼ 110 pc
(diameter) reaches a mass of M ∼ 1.1 × 106 M . In Figure 10 we compare the W49 GMC with the Galactic Center cloud G0.253+0.016 (Longmore et al. 2012; Kauffmann et al. 2013), which has a mass M ∼ 1.3 × 105 M
in a diameter D = 5.6 pc. Also plotted are the most massive clumps of the G305 GMC (Hindson et al. 2010) and
the result of the recent study of the Carina complex by
Preibisch et al. (2012). Unlike W49 and G0.253+0.016,
the clusters in the latter two clouds are already optically
visible. G0.253 and the W49 GMC also differ in other
properties. G0.253 is almost starless (Kauffmann et al.
2013; Rodr´
ıguez & Zapata 2013), appears to be externally confined to a total size of a few pc (Longmore et al.
2013), and is in a more extreme environment in the Central Molecular Zone (CMZ) of the Milky Way that may
difficult further star formation. On the other hand, W49
has already formed a copious amount of stars as seen in
the infrared and radio continuum.
It is seen that GMCs like Carina or W49 have a gas
mass reservoir large enough (∼ 106 M ) and the right
size (radius ∼ 101 − 102 pc) to form stellar clusters with
the typical stellar mass M ∼ 105 M and radius ∼ 10
pc of globular clusters (Portegies Zwart et al. 2010). The
gas-mass reservoir of the central few pc is still enough to
form a YMC with a stellar mass M > 104 M , even if
it decouples completely from the rest of the cloud.
We note that the gas mass is dominated by molecular
gas at all the scales relevant to this discussion. From the
3.5-cm free-free map, and for a range of turn-off frequencies26 from 8 to 230 GHz, the total HII mass is between
MHII ∼ 1700 and 5400 M .
5. GMC AND CLUSTER DISRUPTION

What is the future of the W49 GMC and its star clusters? Will the GMC be disrupted and star formation terminated? Will the resulting star cluster remain bound?
5.1. GMC disruption by feedback

The GMC can be disrupted by a series of feedback processes, from which the most important are protostellar
jets and outflows, HII pressure, and radiation pressure.
In W49A, the luminosity output is heavily dominated by
the most massive stars. Massive, collimated outflows, do
not appear to be important.
In principle, the over-pressurized ionized gas at temperature T ∼ 104 K can clear out the surrounding molecular gas, which is at a temperature from tens to hundreds of Kelvin (Keto 2007). However, W49A is at a
stage at which in spite of its enormous luminosity, several of its HII regions are still confined within ∼ 1000
AU (De Pree et al. 2000, 2004). Various mechanisms
26 The frequency at which an HII region transitions from being
optically thin to thick.

8

a

may be responsible for delaying HII region growth, including confinement by denser molecular gas (De Pree
et al. 1995), molecular infall from the surrounding clump
(Walmsley 1995), and continued accretion through filamentary, partially-ionized accretion flows (Peters et al.
2010a; Galvń-Madrid et al. 2011; Dale et al. 2012).
a
W49A has already a lifetime of several ×105 yr since
the onset of massive star formation27 , given the presence
of hypercompact HII regions and hot molecular cores
(Wilner et al. 2001). After this time, only ∼ 1 % of the
gas mass in the inner ∼ 6 pc around W49N is ionized, the
rest is molecular. The difference is even more dramatic
(∼ 0.1 % of gas is ionized) if the full GMC is considered.
Therefore, it is evident that in this timescale ionization
is not efficient to disrupt the cloud. Even if the few central pc of gas were fully ionized, the main disruptive force
would be radiation pressure rather than ionized-gas pressure. From Figure 2 of Krumholz & Matzner (2009), for
an ionizing photon rate Q ∼ few ×1051 s−1 , and an average particle density n ∼ few ×103 cm−3 , the radiation
pressure would be dominant by a factor ∼ 100.
To estimate the effect of radiation pressure, since most
of the gas is molecular and not ionized, we should consider that the dominant form will be onto dust grains28 .
The GMC is disrupted if the outward radiation pressure
from the central clusters overcomes the cloud’s self gravity, i.e., if the luminosity from the clusters L is greater
than the Eddington luminosity LEdd . Under spherical
symmetry (an assumption only valid to some extent in
the inner ∼ 10 pc of the central clusters):
LEdd (r) = 4πcGM (r)/κR ,

(1)

where M (r) is the dust mass at radius r and κR is the
Rosseland mean dust opacity. We take M (r = 6.5pc) ∼
2.1 × 103 M from our CO and dust measurements29 ,
and κR in the range from 0.1 to 1 cm2 g−1 (Semenov
et al. 2003). Therefore, LEdd (r = 6.5pc) is in the range
from 2.7 × 107 to 2.7 × 108 L . Sievers et al. (1991)
obtained the luminosity of the central clusters by fitting
their spectral energy distribution from the mid infrared
to the millimeter. We rescale their result to a distance of
11.1 kpc (those authors used 14 kpc), to obtain a total
luminosity Lbol ≈ 1.7 × 107 L . The Eddington ratio
Lbol /LEdd is then in the range 0.06 to 0.6. Note that the
radius selection corresponds to the distance of W49N to
W49S, and therefore is roughly the minimum gas mass
that encloses most of the luminosity. If larger radii are
taken, or if the larger mass estimates from taking an
aperture following the central filaments in the CO maps,
or from the 0.4-mm observations are taken, then the Eddington ratio Lbol /LEdd would be even lower. Also, if the
totality of the GMC is considered, it is possible that the
filaments that are seen to radially converge in the central
clusters continue to feed them during the star-formation
timescale. A dynamical study of the filamentary network
will follow this paper, however, even in the static case,
27

The timespan of any previous low-mass star formation, plus
the timespan of the assembly of the GMC are apart.
28 Even photons that are first absorbed by hydrogen atoms and
returned at lower energies will be absorbed by dust (e.g., Krumholz
& Matzner 2009).
29 Note that this is the mass in the dust component.

we can conclude that the current cluster is not able to
entirely clear its own cloud.
Murray et al. (2010) investigated the disruption of
GMCs across a wide range of conditions, from Galactic starburst regions like W49 to Ultra-luminous Infrared Galaxies. They indeed conclude that radiation
pressure is the dominant force in cloud dispersal. For
their W49A case, however, they pick a cluster luminosity
L = 6.4 × 107 L from correcting the free-free luminosity for dust absorption. This higher cluster luminosity
in their model allows it to marginally disperse the GMC.
Although the luminosity derived from the far-IR by Sievers et al. (1991) is more accurate, it is also true that the
cluster may still gain more mass and become more luminous. An appealing idea is that the cluster will continue to gain mass until Lbol /LEdd ≈ 1. Indeed, there
is evidence that star formation in the central clusters is
continuous: part of the stellar population is already visible in the near infrared30 (Conti & Blum 2002; Alves &
Homeier 2003); part of it is barely visible in the nearand mid-IR (Smith et al. 2009a) but their massive stars
are seen as HC and UCHII regions in the radio (e.g., De
Pree et al. 1997, 2000); and finally, there are also hot
molecular cores without associated HII region (Wilner
et al. 2001).
5.2. The fate of the central star cluster
Regarding the future of the star clusters, W49N and
W49S may form more than one YMC, or coalesce into
a single system. They could also dissolve. Observations
and models of YMC dynamics suggest that if the starformation efficiency is low enough (< 50 %) ”rapid”31
gas dispersal leaves unstable star clusters that end up
dissolving (Lada et al. 1984; Bastian & Goodwin 2006).
However, if the gas dispersal timescale is comparable or
larger than the crossing time of the GMC, the resulting
star cluster has more time to adjust to the new gravitational potential and it is easier for it to remain bound
even with star formation efficiencies as low as 10 % (Lada
et al. 1984; Pelupessy & Portegies Zwart 2012).
We argue that, because the star formation efficiency
SF E = M /(M + Mgas ) and the gas dispersal timescale
are large enough, the most likely outcome is that the star
clusters W49N+W49S (or at least part of their stellar
content) will remain bound.
From near-infrared observations, Homeier & Alves
(2005) estimate a stellar mass in frames of length ∼ 5
(∼ 16 pc) of M ∼ 4 × 104 to 7 × 104 M . The total
stellar mass should be somewhat larger when the embedded population not visible in the IR (e.g., the HC
HII regions) is taken into account. The total gas mass
in the same area from the CO and submm continuum
data is Mgas ∼ 4 × 105 to 6 × 105 M . Therefore, within
the W49N+W49S region, the current SF E is > 10 %,
and will likely increase with time until star formation is
terminated.
30 Concentrated ∼ 3 pc east of the Welch ring of HC HIIs, the
”Cluster 1” of Alves & Homeier (2003) likely is the ionizing source
of the extended halo of ionized emission seen in the cm (De Pree
et al. 1997), whose edge matches the shell reported by Peng et al.
(2010).
31 Rapid compared to the crossing time t . For a GMC of size
cr
∼ 10 pc and velocity dispersion ∼ 5 km s−1 , tcr ∼ 2 Myr.

Observations of the most luminous star formation regions in the Galaxy show that their star clusters are still
deeply embedded on large scales even after the formation of subsequent generations of UC HIIs, HC HIIs, and
hot molecular cores, which requires timescales of at least
several ×105 yr from the onset of significant feedback
(Galvń-Madrid et al. 2009; Liu et al. 2011, 2012a,b).
a
Older, optically visible clusters like the one in Carina
(Preibisch et al. 2012) are still surrounded by hundreds
of thousands of solar masses of molecular gas. Therefore,
observations show that gas dispersal is slow compared to
the star formation timescale. It is harder to estimate the
SFR of the whole GMC because although in this paper
we have presented a good account of its gas content, most
of the stellar content is deeply embedded.
The central part of the W49 GMC satisfies other
proposed criteria for bound massive cluster formation.
Bressert et al. (2012) proposed that massive clusters need
to have a escape velocity vesc greater than the sound
speed of ionized hydrogen cHII ≈ 10 km s−1 . The central clusters in W49A satisfy this criterion. For a current star cluster mass Mcl ∼ 7 × 104 M , the minimum
bound-cluster radius is rcl = 2GMcl /c2 ≈ 6 pc. If the
HII
embedded population is taken into account, assuming
Mcl = 105 M , rcl = 8.6 pc. The near-IR clusters32
identified by Alves & Homeier (2003) have individual
radii r < 1 pc. If further contiguous star formation ends
up making a single cluster covering all the central part
of the cloud, from W49S to W49SE, the radius of the
resultant cluster is r < 6 pc, compact enough to satisfy the criterion of Bressert et al. (2012) for boundness.
This, however, may be a very restrictive condition, since
it assumes that even a fully ionized cloud is contained by
gravity and will eventually be used to form stars. Our
observations show that in the W49 GMC the ionized-gas
mass fraction is very small.
Kruijssen (2012) presented a model in which bound
star clusters are formed out of the high density end of
the ISM. This model accounts for possible dissolution
effects like gas expulsion and external tidal forces, for
conditions from typical Milky Way to extragalactic starbursts. Kruijssen (2012) finds that for gas mass surface
densities above Σgas ∼ 1000 M pc−2 , about 70 % of the
formed clusters will remain bound. Our surface density
map (figure 9, bottom right) shows that almost all of the
W49N gas is above this threshold. This further suggests
that the central cluster W49N33 will remain as a bound
one.
6. CONCLUSIONS

The first results from our MUSCLE survey of gas in
the W49 GMC can be summarized as follows:
• The W49 GMC is one of the most massive in
the Galaxy. From multi-scale observations of CO
and isotopologues, we derive a total mass Mgas ∼
1.1 × 106 M within a radius of 60 pc. Around the
most prominent cluster W49N forming at the center of the GMC, within a radius of 6 pc, the total
gas mass is Mgas ∼ 2 × 105 M (masses uncertain
32 “Cluster 1” is associated with W49N, ”cluster 2” with W49S,
and ”cluster 3” with W49SE.
33 We do not have a high-resolution mosaic of W49S.

9

to ±50%). Therefore ∼ 20 % of the gas mass is concentrated in ∼ 0.1 % of the volume. The gas mass
is dominated by the molecular- rather than the ionized phase (only ∼ 1 % of gas is ionized in the inner
region). The W49 GMC has enough mass to form a
young massive cluster (YMC) as massive as a globular cluster with a conservative star formation efficiency. We compare our results with recent studies
of clouds candidate to form YMCs, like the Galactic Center cloud G0.253+0.016 (Longmore et al.
2012), the Carina complex (Preibisch et al. 2012),
and G305 (Hindson et al. 2010).
• The mass of the GMC is distributed in a hierarchical network of filaments that is forming a young
massive cluster (YMC), or a system of YMCs. At
scales < 10 pc, a triple, centrally condensed, filamentary structure peaks toward the ring of hypercompact HII regions in W49N known to host
dozens of deeply embedded (maybe still accreting)
O-type stars. This structure is observed to continue at scales from ∼ 10 to 100 pc through filaments that radially converge toward W49N and
toward its less prominent neighbour W49S. These
large scale filaments are clumpy and could be forming stars at a rate lower than that of the central clusters. This finding suggests that the W49A
starburst most likely formed from global gravitational contraction with localized collapse in a ”hubfilament” geometry.
• Feedback from the central YMCs (with a current
mass Mcl
7 × 104 M ) is still not enough to
disrupt the GMC, but further stellar mass growth
within a factor of 2 could be enough to allow radiation pressure to disrupt the cloud and halt star
formation. There is no evidence on global scales
for significant disruption from photoionization.
• Likely, the resulting stellar content will remain as
a gravitationally bound massive star cluster, or a
small system of bound star clusters.

The authors acknowledge an anonymous referee for
a timely and useful report. The authors are grateful
to Mr. Bing-Gang Ju and Mr. Deng-Rong Lu from
the PMO operations staff. R.G.-M. acknowledges funding from the European Community’s Seventh Framework Programme (/FP7/2007-2013/) under grant agreement No. 229517R. L.F.R. and L.Z. acknowledge the
support of DGAPA, UNAM, and CONACyT (Mexico). T.P. acknowledges financial support through SNF
grant 200020 137896. R.G.-M. and E.K. acknowledge
the hospitality of the Aspen Center for Physics, which
is supported by the National Science Foundation Grant
No. PHY-1066293. This research made use of APLpy,
an open-source plotting package for Python hosted at
http://aplpy.github.com. The authors acknowledge the
anonymous referee for an encouraging and useful report.
R. G.-M. thanks Quang Nguyen Luong and Diederik
Kruijssen for comments on a draft of the paper. J.E.P
has received funding from the European Community’s

10

a

Seventh Framework Programme (/FP7/2007-2013/) under grant agreement No 229517 and from the SNF (Swiss
National Science Foundation) Sinergia Project.
REFERENCES
Aguirre, J. E., Ginsburg, A. G., Dunham, M. K., et al. 2011,
ApJS, 192, 4
Alves, J., & Homeier, N. 2003, ApJ, 589, L45
Baobab Liu, H., Ho, P. T. P., Zhang, Q., et al. 2010, ApJ, 722,
262
Bastian, N., & Goodwin, S. P. 2006, MNRAS, 369, L9
Belloche, A., M¨ller, H. S. P., Menten, K. M., Schilke, P., &
u
Comito, C. 2013, ArXiv e-prints
Beltrń, M. T., Cesaroni, R., Zhang, Q., et al. 2011, A&A, 532,
a
A91
Beuther, H., Linz, H., Henning, T., et al. 2011, A&A, 531, A26
Bonnell, I. A., Bate, M. R., Clarke, C. J., & Pringle, J. E. 2001,
MNRAS, 323, 785
Bressert, E., Ginsburg, A., Bally, J., et al. 2012, ApJ, 758, L28
Buckley, H. D., & Ward-Thompson, D. 1996, MNRAS, 281, 294
Cesaroni, R., Hofner, P., Araya, E., & Kurtz, S. 2010, A&A, 509,
A50
Col´ P., V´zquez-Semadeni, E., & G´mez, G. C. 2013, MNRAS
ın,
a
o
Conti, P. S., & Blum, R. D. 2002, ApJ, 564, 827
Curran, R. L., & Chrysostomou, A. 2007, MNRAS, 382, 699
Dale, J. E., Ercolano, B., & Bonnell, I. A. 2012, MNRAS, 424, 377
De Pree, C. G., Mehringer, D. M., & Goss, W. M. 1997, ApJ,
482, 307
De Pree, C. G., Rodriguez, L. F., & Goss, W. M. 1995, Revista
Mexicana de Astronomia y Astrofisica, 31, 39
De Pree, C. G., Wilner, D. J., Goss, W. M., Welch, W. J., &
McGrath, E. 2000, ApJ, 540, 308
De Pree, C. G., Wilner, D. J., Mercer, A. J., et al. 2004, ApJ,
600, 286
Di Francesco, J., Johnstone, D., Kirk, H., MacKenzie, T., &
Ledwosinska, E. 2008, ApJS, 175, 277
Dickel, H. R., & Goss, W. M. 1990, ApJ, 351, 189
Duarte-Cabral, A., Dobbs, C. L., Peretto, N., & Fuller, G. A.
2011, A&A, 528, A50
Fukui, Y., Ohama, A., Hanaoka, N., et al. 2013, ArXiv e-prints
Galvń-Madrid, R., Keto, E., Zhang, Q., et al. 2009, ApJ, 706,
a
1036
Galvń-Madrid, R., Peters, T., Keto, E. R., et al. 2011, MNRAS,
a
416, 1033
Galvń-Madrid, R., Zhang, Q., Keto, E., et al. 2010, ApJ, 725, 17
a
Ginsburg, A., Glenn, J., Rosolowsky, E., et al. 2013, ApJS, 208,
14
Gooch, R. 1996, in Astronomical Society of the Pacific
Conference Series, Vol. 101, Astronomical Data Analysis
Software and Systems V, ed. G. H. Jacoby & J. Barnes, 80
Gutermuth, R. A., Myers, P. C., Megeath, S. T., et al. 2008, ApJ,
674, 336
Gwinn, C. R., Moran, J. M., & Reid, M. J. 1992, ApJ, 393, 149
Hindson, L., Thompson, M. A., Urquhart, J. S., Clark, J. S., &
Davies, B. 2010, MNRAS, 408, 1438
Ho, L. C. 1997, in Revista Mexicana de Astronomia y Astrofisica
Conference Series, Vol. 6, Revista Mexicana de Astronomia y
Astrofisica Conference Series, ed. J. Franco, R. Terlevich, &
A. Serrano, 5
Ho, P. T. P., & Haschick, A. D. 1981, ApJ, 248, 622
Ho, P. T. P., Moran, J. M., & Lo, K. Y. 2004, ApJ, 616, L1
Homeier, N. L., & Alves, J. 2005, A&A, 430, 481
Inoue, T., & Fukui, Y. 2013, ApJ, 774, L31
Johnson, K. E., Kobulnicky, H. A., Massey, P., & Conti, P. S.
2001, ApJ, 559, 864
Kauffmann, J., Pillai, T., & Zhang, Q. 2013, ApJ, 765, L35
Keto, E. 2007, ApJ, 666, 976
Keto, E., Ho, L. C., & Lo, K.-Y. 2005, ApJ, 635, 1062
Keto, E., Zhang, Q., & Kurtz, S. 2008, ApJ, 672, 423
Keto, E. R., Lattanzio, J. C., & Monaghan, J. J. 1991, ApJ, 383,
639
Kirk, H., Myers, P. C., Bourke, T. L., et al. 2013, ApJ, 766, 115
Kruijssen, J. M. D. 2012, MNRAS, 426, 3008
Krumholz, M. R., Klein, R. I., & McKee, C. F. 2011, ApJ, 740, 74
Krumholz, M. R., & Matzner, C. D. 2009, ApJ, 703, 1352

Kurtz, S., Cesaroni, R., Churchwell, E., Hofner, P., & Walmsley,
C. M. 2000, Protostars and Planets IV, 299
Kurtz, S., Churchwell, E., & Wood, D. O. S. 1994, ApJS, 91, 659
Lada, C. J., Margulis, M., & Dearborn, D. 1984, ApJ, 285, 141
Liu, H. B., Ho, P. T. P., Wright, M. C. H., et al. 2013, ApJ, 770,
44
Liu, H. B., Jimńez-Serra, I., Ho, P. T. P., et al. 2012a, ApJ, 756,
e
10
Liu, H. B., Quintana-Lacaci, G., Wang, K., et al. 2012b, ApJ,
745, 61
Liu, H. B., Zhang, Q., & Ho, P. T. P. 2011, ApJ, 729, 100
Longmore, S. N., Rathborne, J., Bastian, N., et al. 2012, ApJ,
746, 117
Longmore, S. N., Kruijssen, J. M. D., Bally, J., et al. 2013,
MNRAS, 433, L15
Lovas, F. J. 2004, Journal of Physical and Chemical Reference
Data, 33, 177
Mangum, J. G., Emerson, D. T., & Greisen, E. W. 2007, A&A,
474, 679
Matthews, H., Kirk, H., Johnstone, D., et al. 2009, AJ, 138, 1380
Meurer, G. R., Heckman, T. M., Leitherer, C., et al. 1995, AJ,
110, 2665
Miyawaki, R., Hayashi, M., & Hasegawa, T. 1986, ApJ, 305, 353
—. 2009, PASJ, 61, 39
Molinari, S., Swinyard, B., Bally, J., et al. 2010, A&A, 518, L100
M¨ller, H. S. P., Schl¨der, F., Stutzki, J., & Winnewisser, G.
u
o
2005, Journal of Molecular Structure, 742, 215
Murray, N., Quataert, E., & Thompson, T. A. 2010, ApJ, 709,
191
Myers, P. C. 2009, ApJ, 700, 1609
Nagy, Z., van der Tak, F. F. S., Fuller, G. A., Spaans, M., &
Plume, R. 2012, A&A, 542, A6
Nguyen Luong, Q., Motte, F., Schuller, F., et al. 2011, A&A, 529,
A41
Nguyen Luong, Q., Motte, F., Carlhoff, P., et al. 2013, ArXiv
e-prints
Nummelin, A., Bergman, P., Hjalmarson, A., et al. 1998, ApJS,
117, 427
Ossenkopf, V., & Henning, T. 1994, A&A, 291, 943
Pelupessy, F. I., & Portegies Zwart, S. 2012, MNRAS, 420, 1503
Peng, T.-C., et al., & et al. 2013, submitted
Peng, T.-C., Wyrowski, F., van der Tak, F. F. S., Menten, K. M.,
& Walmsley, C. M. 2010, A&A, 520, A84
Peters, T., Banerjee, R., Klessen, R. S., & Mac Low, M.-M. 2011,
ApJ, 729, 72
Peters, T., Banerjee, R., Klessen, R. S., et al. 2010a, ApJ, 711,
1017
Peters, T., Mac Low, M.-M., Banerjee, R., Klessen, R. S., &
Dullemond, C. P. 2010b, ApJ, 719, 831
Pineda, J. E., Goodman, A. A., Arce, H. G., et al. 2011, ApJ,
739, L2
Pineda, J. E., Rosolowsky, E. W., & Goodman, A. A. 2009, ApJ,
699, L134
Pon, A., Toal´, J. A., Johnstone, D., et al. 2012, ApJ, 756, 145
a
Portegies Zwart, S. F., McMillan, S. L. W., & Gieles, M. 2010,
ARA&A, 48, 431
Preibisch, T., Roccatagliata, V., Gaczkowski, B., & Ratzka, T.
2012, A&A, 541, A132
Roccatagliata, V., Preibisch, T., Ratzka, T., & Gaczkowski, B.
2013, A&A, 554, A6
Rodr´
ıguez, L. F., & Zapata, L. A. 2013, ApJ, 767, L13
Santangelo, G., Testi, L., Gregorini, L., et al. 2009, A&A, 501, 495
Sault, R. J., Teuben, P. J., & Wright, M. C. H. 1995, in
Astronomical Society of the Pacific Conference Series, Vol. 77,
Astronomical Data Analysis Software and Systems IV, ed.
R. A. Shaw, H. E. Payne, & J. J. E. Hayes, 433
Semenov, D., Henning, T., Helling, C., Ilgner, M., & Sedlmayr, E.
2003, A&A, 410, 611
Serabyn, E., Guesten, R., & Schulz, A. 1993, ApJ, 413, 571
Shan, W. L., Yang, J., Shi, S. C., et al. 2010, in Twenty-First
International Symposium on Space Terahertz Technology, 136
Sievers, A. W., Mezger, P. G., Bordeon, M. A., et al. 1991, A&A,
251, 231
Simon, R., Jackson, J. M., Clemens, D. P., Bania, T. M., &
Heyer, M. H. 2001, ApJ, 551, 747
Smith, N., Whitney, B. A., Conti, P. S., de Pree, C. G., &
Jackson, J. M. 2009a, MNRAS, 399, 952

Smith, R. J., Longmore, S., & Bonnell, I. 2009b, MNRAS, 400,
1775
Snell, R. L., Scoville, N. Z., Sanders, D. B., & Erickson, N. R.
1984, ApJ, 284, 176
Takahashi, S., Ho, P. T. P., Teixeira, P. S., Zapata, L. A., & Su,
Y.-N. 2013, ApJ, 763, 57
Toal´, J. A., V´zquez-Semadeni, E., & G´mez, G. C. 2012, ApJ,
a
a
o
744, 190
Turner, J. L. 2009, Extreme Star Formation, ed. H. A. Thronson,
M. Stiavelli, & A. Tielens, 215
V´zquez-Semadeni, E., G´mez, G. C., Jappsen, A.-K.,
a
o
Ballesteros-Paredes, J., & Klessen, R. S. 2009, ApJ, 707, 1023

11

Walmsley, M. 1995, in Revista Mexicana de Astronomia y
Astroﬁsica, vol. 27, Vol. 1, Revista Mexicana de Astronomia y
Astroﬁsica Conference Series, ed. S. Lizano & J. M. Torrelles,
137
Ward-Thompson, D., & Robson, E. I. 1990, MNRAS, 244, 458
Wei, L. H., Keto, E., & Ho, L. C. 2012, ApJ, 750, 136
Welch, W. J., Dreher, J. W., Jackson, J. M., Terebey, S., &
Vogel, S. N. 1987, Science, 238, 1550
Whitmore, B. C., Schweizer, F., Leitherer, C., Borne, K., &
Robert, C. 1993, AJ, 106, 1354
Wilner, D. J., De Pree, C. G., Welch, W. J., & Goss, W. M. 2001,
ApJ, 550, L81
Wilson, T. L., & Rood, R. 1994, ARA&A, 32, 191
Zhang, B., Reid, M. J., Menten, K. M., et al. 2013, to appear in
ApJ

12

a

APPENDIX
A: DETAILS OF THE SMA MOSAICS

The SMA observations make use of the 4 available array configurations to cover the baseline range from 8 to 480
kλ (Fig. A1, left). Interferometric observations of massive star formation regions (and in general, of fields with
complicated structure at multiple scales) are typically limited by the dynamic range produced by the discrete u, v
coverage rather than by integration time as in point-source observations. We have corroborated in our data set, as well
as in other studies over the past years (e.g., Galvń-Madrid et al. 2010) that co-adding SMA compact and extended
a
configurations helps to recover all the flux up to scales of ∼ 10 (comparable to the core scales of ∼ 0.1 pc at typical
kpc distances) while preserving subarcsecond angular resolution. However, since the ISM is hierarchical, cores are
themselves embedded in pc scale clumps which may themselves be embedded on larger-scale structures, as shown in
this paper. To map these structures, combined single-dish and inteferometer mapping is needed. ALMA is supposed
to resolve these issues by combining mapping from two different arrays and single dishes. The data presented here
covers a larger range of angular scales (plus the single dish) than even what has been offered for ALMA Cycle 1.
The right panel of Figure A1 shows the primary beam response of the 11-pointing mosaic. The images and measurements made in this paper have been divided by this reponse to correct the flux scale. Most of the interferometer
maps have most of the flux in the central part of the mosaic. However, the CO (and isotopolgues) and dust continuum
maps suffer from missing flux even in the all-configuration mosaics, so we have combined them with single-dish data.
The SMA continuum map has been combined with the BOLOCAM GPS survey map (Ginsburg et al. 2013). The
13
CO and C18 O 2–1 maps have been combined with IRAM 30m maps presented in Peng et al. (2013). The combining
procedure is described in Liu et al. (2012b, 2013). It is based on creating a visibility data set from the single-dish
maps that overlaps in u, v range with the interferometer data. The concatenated data is then inverted back to the
image space and jointly cleaned.
B: PMO VELOCITY-INTEGRATED MAPS

In this Appendix we show the velocity-integrated (moment 0) PMO maps of the lines listed in Table 2, except for
the CO maps that are used in the body of the text.
C: SMA VELOCITY-INTEGRATED MOSAICS

In this Appendix we show the velocity-integrated SMA maps from all the identified lines listed in Table 5, except for
the 13 CO and C18 O which are used in the main body of the text. A detailed analysis of these maps will be presented
in a forthcoming paper.
D: CALCULATIONS OF THE MASS OF THE GMC

We calculate the mass of the W49 GMC at all scales using four data sets: the combined SMA+IRAM 30m 13 CO
and C18 O 2–1 maps of the inner ∼ 10 pc (Fig. 9), the combined SMA+BOLOCAM GPS 1 mm dust continuum map
also in the inner ∼ 10 pc (Fig. 5), the archive SCUBA 0.4-mm map (Fig. 6), and the CO and 13 CO 1–0 PMO maps
covering the full GMC up to ∼ 110 pc scales (Fig. 9). Below we describe the methodology.
CO 1–0
If we assume that the 13 CO emission is optically thin, the 12 CO opacity (τ12 CO ) in a given voxel of data (a positionposition-velocity cell) can be calculated from the intensity (brightness) ratio of the lines (Snell et al. 1984):
TB,12 CO
1 − exp(−τ12 CO )
=
,
TB,13 CO
1 − exp(−τ12 CO /χ13 C )

(D1)

where χ13 C is the relative abundance of 12 C to 13 C. We use χ13 C = 65 using the fit of Wilson & Rood (1994) for a
Galactocentric distance of 7.6 kpc for W49, derived using its parallax distance to the Sun of 11.1 kpc (Zhang et al.
2013).
In the voxels where there is a 12 CO detection but no 13 CO we use the optically thin approximation to the column
density of 12 CO [cm−2 ]:
Nthin,12 CO = 4.31 × 1013

Tex + 0.9
TB,12 CO ∆v,
exp(−5.5/Tex )

(D2)

where temperatures are in K and velocity widths in km s−1 .
In the voxels with 13 CO emission we calculate the optically-thick column density with:
Nthick,12 CO =

τ12 CO
Nthin,12 CO .
1 − exp(−τ12 CO )

(D3)

Once the column density is known for each position-position-velocity voxel, we integrate them to obtain the mass
surface density maps shown in Section 4. We note that in principle the same procedure can be used with the C18 O to


13

12
CO ratio. However, the PMO C18 O maps are not sensitive enough to detect emission except for the very center of
the cloud (see Section 3.1).
Two small corrections have been applied to the straightforward procedure outlined above. First, there is a dip in
the spectra between 14 and 18 km s−1 that appears to be caused by an extended, foreground cloud (Peng et al. 2013).
We have interpolated the measurements within this range. The corrected mass is < 1% larger. Second, a few voxels
close to the map center and around ∼ 7 km s−1 have very large optical depths (i.e., the 12 CO and 13 CO maps have
the same intensity within the uncertainty).
For a cloud that is fully molecular, a hydrogen to helium number ratio of 10, and an abundance of H2 to 12 CO
χCO = 104 , the gas mass surface density Σ is then obtained from:

Σ = µmH χCO N12 CO ,

(D4)

where [Σ] = g cm−2 for [N12 CO ] = cm−2 , µ = 2.8, and mH = 1.6733 × 10−24 g.
CO 2–1
On the scales of W49N (∼ 10 pc) we follow a similar procedure to the one outlined above, but with the 13 CO and
18
C O 2–1 maps at ∼ 2 angular resolution combined from the SMA mosaics and IRAM 30m data. In this case, we
have that:
Nthin,13 CO = 1.17 × 1013 exp(5.3/Tex )

Tex + 0.9
TB,13 CO ∆v.
exp(−10.6/Tex )

(D5)

Th rest of the procedure is as in the larger-scale maps. The relative abundance of 13 CO to C18 O (7.5) is the ratio of
the abundance of 16 O to 18 O (χ18 O = 484.3) to the abundance of 12 C to 13 C (χ13 C = 65) (Wilson & Rood 1994). The
maximum optical depth τC18 O solved in our calculations is 20. As with the larger-scale GMC, we have interpolated
the spectra within the velocity range from 14 and 18 km s−1 .
Millimeter continuum
Dust grains can be assumed to emit as a gray body at temperature Td . At millimeter and submillemeter wavelenghts
their emission is optically thin, so the total mass in dust grains is given by
Md = 3.25 × 106 d2

exp(0.048ν/Td ) − 1 Sν
,
ν3
κν

(D6)

where the units of mass [Md ], distance [d], frequency [ν], dust temperature [Td ], flux [Sν ], and dust absorption coefficient
[κν ] are respectively M , kpc, GHz, K, Jy, and cm2 g−1 .
The total molecular-gass mass is then Mgas = 100Md . Often the Rayleigh-Jeans regime (hν
kTd ) is assumed, but
this is not the case for Td ∼ 30 K and the frequencies we deal with, specially at 678 GHz (0.4 mm) where hν/kTd ≈ 1.
Therefore we prefer to use the full Planck function.
Uncertainties
Many sources of uncertainty make the estimation of masses derived from CO ratios and dust continuum emission
accurate to a factor of 2, at most. For the CO-ratio measurements, we handle uncertainties on the isotopologue
abundances and the excitation temperature, and do not consider uncertainties in the CO to H2 abundance factor,
which is usually set in the literature to χCO = 104 . For the dust measurements, we handle uncertainties in the dust
absorption coefficient and the dust temperature, but again leave the gas-to-dust mass conversion factor at 100.
The uncertainty in the mass determinations from CO isotopologue abundances is estimated from the fits of Wilson
& Rood (1994). The abundance of 12 C with respect to 13 C is χ13 C = 65 ± 27, and the abundance of 16 O to 18 O
is χ18 O = 484 ± 172. For the optically-thin and the non-saturated optically thick voxels, uncertainty in the mass
propagates as the uncertainty in abundance. However, this is not the case for the saturated voxels in which we set τ
to a maximum value τmax . τmax is defined as the value at which the line brightness ratio equals 1 + σ, where σ is the
rms noise in the maps, i.e., larger optical depths are not distinguishable, within the noise. We have found empirically
how this affect our calculations by re-running them within the allowed range of abundance values. In the large-scale
PMO maps, very few pixels are affected by saturation in all cases. Therefore the masses scale almost linearly with the
used abundance. In the smaller-scale SMA maps, several voxels are saturated mainly in the central parts of the map,
therefore larger abundances permit increasingly large τmax and larger mass estimates. The mass uncertainty due to
the above mentioned effect grows at smaller radii, up to +130% and −57% within r = 0.5 pc.
If the excitation temperature Tex is not position dependent, then its effect on the uncertainty of the mass estimate
is almost linear and affects equally all the radial measurements. We pick Tex = 30 K. Reasonable Tex values are in the
range from 20 to 40 K. Therefore we estimate a respective mass uncertainty of ±26%. We also explore the effect of a
possible radial dependence Tex ∝ r−1.5 , with Tex (r = 0) = 50 K, and Tex (r = 5pc) = 20 K (i.e., a warm W49N). The
results show that the mass in the inner 5 pc would be systematically underestimated by ∼ 54% close to the center and
∼ 27% at r = 5 pc, and then at larger radii the mass difference would converge to ∼ 3% at r = 50 pc. However, we
do not include the effects of this hypothetical temperature profile in our error budget.

14

a

Propagating the uncertainties from abundances and excitation temperatures gives a total uncertainty of ∼ ±50%
for the PMO masses. For the SMA+IRAM 30m masses, typical values are +80% and −50%, but uncertainties are
larger at smaller radii. The error bars in Figure 10 include the individual contributions from abundance uncertainty
and from shifts in Tex .
Uncertainties in mass determinations from dust are as large or larger than from CO. The dust absorption coefficient
κν , dependent on the properties of dust grains, is the main source of uncertainty. We take as possible values for κν
those corresponding to coagulated dust grains without significant ice mantles from the calculations of Ossenkopf &
Henning (1994). The dispersion of these values is smaller at 0.4 mm than at 1 mm. We estimate the mass uncertainty
from κν to be ± ∼ 20% at the former wavelength, and ± ∼ 65% at the latter. Possible variations from the adopted
dust temperature Td = 30 K also contribute to the total uncertainty in mass. From the gray body fits of Sievers
et al. (1991) and Ward-Thompson & Robson (1990), Td is restricted to the range between 20 and 50 K. Therefore the
uncertainty from temperature in the mass determination is ∼ ±50%. Again, we do not consider further uncertainties
from the dust-to-gas mass conversion factor, which we fix at 100. Therefore, we estimate a total mass uncertainty of
±55% at 0.4 mm, and ±90% at 1 mm, including uncertainties from the free-free subtraction at the latter wavelength.


15

Figure 1. PMO velocity-integrated (moment 0, from −20 to 30 km s−1 ) CO intensity maps of the entire W49 GMC. Units are K km
s−1 . The HPBW is 58 . The scale bar of length 10 is equivalent to 32.3 pc. (a) CO 1–0. The rms in the cleaner zones is ∼ 15 K km
s−1 , and is dominated by faint, extended emission. (b) 13 CO 1–0. rms ∼ 6 K km s−1 . (c) C18 O 1–0. rms ∼ 6 K km s−1 . The main
subcomponents of W49A: W49N, W49S, and W49SW are marked in the bottom panel.

16

a

Figure 2. PMO CO spectra and Gaussian fits at the position of the emission peak, α(J2000) = 19h 10m 14.2s , δ(J2000) = 9◦ 6 23 .
The 12 CO 1–0 spectrum (top) was not fit because it is very complex, whereas the 13 CO 1–0 and C18 O 1–0 spectra (middle and bottom,
respectively) are well fitted with a sum of three Gaussians.


17

Figure 3. PMO spectra of carbon molecules (other than CO and its isotopologues) and of the H41α RL. at the position of the emission
peak, α(J2000) = 19h 10m 14.2s , δ(J2000) = 9◦ 6 23 . Gaussian ﬁts are plotted as gray dashed lines.

18

a

Figure 4.
PMO spectra of non-carbon molecules and Gaussian ﬁts at the position of the emission peak, α(J2000)
19h 10m 14.2s , δ(J2000) = 9◦ 6 23 .

=


19

Figure 5. Top: BOLOCAM GPS 1.1-mm continuum map (color) overlaid on contours at ×0.2, 0.4, 0.6, and 0.8 the peak primary-beam
response of the SMA mosaic on the main cluster (W49N). The second brightest peak 2.5’ southeast of W49N is the secondary cluster
usually referred to as W49 South (W49S). HPBW= 33 × 33 . The rms noise is ∼ 40 mJy. The peak intensity is 31.7 Jy beam−1 . Bottom:
SMA mosaic combined with the BOLOCAM single-dish map. The HPBW of the jointly cleaned image is 2.50 × 2.26 , PA= 75.7◦ . The
rms noise is ∼ 5 mJy. The peak intensity is 3.30 Jy beam−1 .

20

a

Figure 6. JCMT-SCUBA map of the central part of the W49 GMC at 0.4 mm (678 GHz). The most prominent emission comes from
the W49N cluster. W49S is also seen 2.5’ southeast of W49N, as well as the ﬁlamentary structures connecting them and to the rest of the
GMC. HPBW= 8.1 . The rms noise is ∼ 3 Jy. The peak intensity is 487 Jy beam−1 .


21

Figure 7. Full-sideband SMA spectra toward W49A, averaged over an area interior to the 5σ contour of the 1.3-mm map shown in
Figure 5, bottom. All the spectral features above an intensity of 50 mJy beam−1 have been identiﬁed as a single or combination of spectral
lines, and are marked by vertical dashed lines. The respective molecule or atom is labeled, except in cases of blended lines. The complete
information is on Table 5.

22

a

Figure 8. Velocity-integrated (moment 0) maps of CO 2–1 isotopologues in W49A (central ∼ 10 pc). These images show the SMA
(subcompact, compact, extended and very extended conﬁgurations) maps combined with IRAM 30m observations. The color scale shows
the CO emission in Jy beam−1 km s−1 . Contours show the 8.5 GHz (3.6 cm) continuum tracing free-free emission from the embedded HII
regions. VLA HPBW=0.80 ×0.78 , PA=−63.1◦ . Top left: Zoomed-out 13 CO 2–1 covering the entire SMA mosaic. HPBW=2.17 ×2.13 ,
PA=−72.1◦ . rms∼ 4.8 Jy beam−1 km s−1 . The massive cluster detected in the NIR by Homeier & Alves (2005) is marked by the red
star. Top right: Zoomed-in 13 CO 2–1 around the central ”Welch” ring of UC and HC HII regions. Bottom left: Zoomed-out C18 O 2–1.
HPBW=2.01 × 1.71 , PA=68.3◦ . rms∼ 0.9 Jy beam−1 km s−1 . Bottom right: Zoomed-in C18 O 2–1.


23

Figure 9. Mass surface density Σ maps obtained from CO-isotopologue line ratios, as described in Appendix C. Units are M pc−2 .
The top frame shows the zoomed-out Σ measurement from the PMO CO and 13 CO 1–0 maps. HPBW=58 . The bottom frame shows the
zoomed-in Σ measurement from the SMA mosaics combined with IRAM 30 maps of 13 CO and C18 O 2–1, covering the area marked by a
black square in the top frame. The contour in the bottom frame corresponds to 1000 M pc−2 . HPBW=2.2 × 2.2 .

24

a

Figure 10. Mass vs radius for Galactic molecular clouds that may form (G0.253) or are indeed forming (the rest of plotted clouds)
massive clusters (M > 104 M in stellar mass). The ﬁlled red circles and blue squares show the total mass in the W49 GMC as a function
of radius around the Welch ring from the observations presented in this paper. The black symbols are measurements compiled from the
literature (only one value of mass and radius is given): the Galactic Center cloud G0.253 (Longmore et al. 2012), the Carina complex
GMC (Preibisch et al. 2012), and the most massive clumps in G305 (Hindson et al. 2010). The typical regimes of stellar mass and radius
for Galactic young massive clusters (YMCs) and globular clusters (GCs) are marked with boxes (Portegies Zwart et al. 2010). Error bars
include systematic uncertainties from element/isotopologue abundances and excitation temperature (see Appendix D).


25

Figure A1. Left: u, v coverage in kilolambdas (kλ) of the concatenated SMA data set from the four array conﬁgurations. Right: Primarybeam response of the 11-pointing SMA mosaic. The images presented in this paper have been divided by this response to correct their ﬂux
scale.

26

a

Figure B1. PMO velocity-integrated (from −12 to 24 km s−1 ) intensity maps of carbon molecules (beside CO and its isotopologues)
with clear detections, plus the H41α recombination line. Units are K km s−1 . The HPBW is 58 . 10 are equivalent to 32.3 pc. The
zoomed areas are diﬀerent because the map sizes are diﬀerent. The rms noises in the shown areas are 5 to 6 K km s−1 in all cases.


27

Figure B2. PMO velocity-integrated (from −12 to 24 km s−1 ) intensity maps of non-carbon molecules with clear detections. Units are
K km s−1 . The HPBW is 58 . 10 are equivalent to 32.3 pc. The zoomed areas are diﬀerent because the map sizes are diﬀerent. The
rms noises in the shown areas are 4 to 5 K km s−1 in all cases.

28

a

Figure C1. (a) Velocity-integrated (moment 0) SMA mosaics, a: The color scale shows the velocity-integrated intensity in Jy beam−1 km
s−1 using a linear stretching. The synthesized beam of these maps varies smoothly from 1.02 × 0.69 , PA= 76.8◦ at the lowest frequency
(36.4 K per Jy beam−1 ) to 0.98 × 0.64 , PA= 74.1◦ at the highest (31.5 K per Jy beam−1 ). White contours show the 3.6-cm free-free
continuum at −4, 4, 16, 64 and 256 × 0.5 mJy beam−1 , VLA HPBW=0.80 × 0.78 , PA=−63.1◦ .


Figure C1. (b) Velocity-integrated (moment 0) SMA mosaics, see a for description.

29

30

a

Figure C1. (c) Velocity-integrated (moment 0) SMA mosaics, see a for description.


Figure C1. (d) Velocity-integrated (moment 0) SMA mosaics, see a for description.

31

32

a

Figure C1. (e) Velocity-integrated (moment 0) SMA mosaics, see a for description.


Figure C1. (f) Velocity-integrated (moment 0) SMA mosaics, see a for description.

33

34

a

Figure C1. (g) Velocity-integrated (moment 0) SMA mosaics, see a for description.


Figure C1. (h) Velocity-integrated (moment 0) SMA mosaics, see a for description.

35

36

a
Table 1
SMA observations
Array

Epoch

Flux
Calibrator

Phase
Calibrator

Bandpass
Calibrator

τ225GHz

Subcompact
Compact
Extended
Very extended

15 Aug 2010
19 Jun 2010
10 Sep 2010
15 Jul 2010

Callisto
Neptune
Callisto
Callisto

1751+096
1751+096
1751+096
1751+096

3C454.3
3C279
3C454.3
3C454.3

0.05 to 0.08
0.08
0.08
0.08

Table 2
Lines detected with PMOa
Molecule

Transition

ν0
[GHz]

EU
[K]

Integrated flux
[Jy km s−1 ]

12 CO

C18 O

1−0
1−0
1−0

115.2712
110.2013
109.7821

5.53
5.29
5.27

5.005268 × 106
7.35425 × 105
3.0754 × 104

HCN
H13 CN
HCO+
H13 CO+
CS
H

1−0
1−0
1−0
1−0
2−1
41α

88.6316
86.3399
89.1885
86.7542
97.9809
92.0344

4.25
4.14
4.28
4.16
7.05
–

5.5959 × 104b
1.445 × 103
5.4033 × 104b
1.121 × 103
5.4608 × 104b
1.262 × 103

N2 H+
SiO
SO
SO2

1−0
2−1
2(2) − 1(1)
8(1, 7) − 8(0, 8)

93.1734
86.8469
86.0939
83.6880

4.47
6.25
19.31
36.71

6.273 × 103
1.814 × 103
7.81 × 102
5.25 × 102

13 CO

a Lines detected with the PMO 14m telescope. The first four columns

indicate, respectively, the molecule or atom, the transition, the rest frequency, and the upper-level energy of the transition. The fifth column
lists the velocity-integrated flux of the maps in Jy beam−1 km s−1 .
The intensity conversion factor is 0.027 K per Jy beam−1 at the rest
frequency of the CO 1–0. The noise in the moment 0 maps ranges from
2 to 10 K km s−1 .
b From the 12 C to 13 C line ratios, the average optical depths of the
main isotopologues are τHCN = 1.4, τHCO+ = 0.8, and τCS < 0.1.

Table 3
Line Fitting at Peak Position of PMO dataa
Molecule

Transition

or Atom

Component 1

Component 2

Component 3

V
[km/s]

A
[K]

FWHM
[km/s]

V
[km/s]

A
[K]

FWHM
[km/s]

V
[km/s]

A
[K]

FWHM
[km/s]

C18 O

1−0
1−0
1−0

–
3.9 ± 0.2
4.0 ± 0.4

–
9.3 ± 0.2
0.9 ± 0.1

–
8.5 ± 0.5
8.7 ± 0.9

–
12.0 ± 0.2
12.3 ± 0.3

–
11.8 ± 0.3
1.3 ± 0.2

–
5.9 ± 0.6
6.3 ± 0.7

–
16.9 ± 0.5
16.4 ± 0.3

–
1.7 ± 0.5
0.3 ± 0.2

–
4.1 ± 0.9
1.0 ± 0.8

HCN
H13 CN
HCO+
H13 CO+
CS
H

1−0
1−0
1−0
1−0
2−1
41α

2.5 ± 0.3
−2.2 ± 0.9
2.0 ± 0.3
2.9 ± 0.5
3.2 ± 0.2
12.2 ± 0.6

3.4 ± 0.2
0.3 ± 0.2
5.1 ± 0.3
0.3 ± 0.2
3.6 ± 0.2
0.4 ± 0.1

13.8 ± 2.5
8.3 ± 3.4
7.7 ± 2.0
5.1 ± 2.7
7.9 ± 1.9
30.2 ± 4.2

11.8 ± 0.3
7.6 ± 0.9
10.8 ± 0.3
10.3 ± 1.0
11.7 ± 0.2
–

2.3 ± 0.4
0.3 ± 0.1
3.6 ± 0.3
0.4 ± 0.1
4.8 ± 0.2
–

4.0 ± 1.8
9.3 ± 3.5
4.8 ± 1.8
13.2 ± 3.9b
7.8 ± 1.7
–

18.0 ± 0.9
–
–
–
–
–

0.8 ± 0.2
–
–
–
–
–

10.9 ± 3.7
–
–
–
–
–

N2 H+
SiO
SO
SO2

1−0
2−1
2(2) − 1(1)
8(1, 7) − 8(0, 8)

3.0 ± 0.7
6.9 ± 0.4
9.1 ± 0.7
8.4 ± 0.8

0.4 ± 0.2
0.7 ± 0.1
0.4 ± 0.2
0.3 ± 0.2

8.6 ± 2.9
12.8 ± 2.6
10.7 ± 3.4
13.6 ± 3.9

12.6 ± 0.7
–
–
–

0.6 ± 0.1
–
–
–

10.6 ± 3.0
–
–
–

–
–
–
–

–
–
–
–

–
–
–
–

12 CO
13 CO

a Line parameters from Gaussian fitting at the peak of integrated emission in the PMO maps: α(J2000) = 19h 10m 14.2s , δ(J2000) = 9◦ 6 23 . The
first two columns indicate, respectively, the molecule or atom and the transition. The 12 CO spectrum is too complex to be fitted by 3 or less Gaussians.
Most of the lines are better described by two Gaussian components, whereas a few of them are single-peaked (e.g., H41α).
b The second peak of H13 CO+ 1–0 appears quite wide, possibly due to a baseline difference at both extremes of the line emission.


37

Table 4
Mass structure of W49N from SMA+BOLOCAM
Intensity threshold
I > [Jy beam−1 ]

Surface density threshold
Σ > [g cm−2 ]

Flux from dust
[Jy]

H2 mass
[M ]

0.025
0.050
0.100
0.200
0.400
0.800
1.600

2.1 ± 1.4
4.0 ± 2.7
8.1 ± 5.4
16.3 ± 10.8
32.6 ± 21.6
65.0 ± 43.3
130.0 ± 86.7

17.017
12.745
9.216
7.462
6.267
4.253
2.239

1.484 × 105 ± 9.89 × 104
1.111 × 105 ± 7.41 × 104
8.03 × 104 ± 5.36 × 104
6.51 × 104 ± 4.33 × 104
5.47 × 104 ± 3.64 × 104
3.71 × 104 ± 2.47 × 104
1.95 × 104 ± 1.30 × 104

2
−1
a The masses and threshold mass surface densities correspond to κ
(per unit
1mm = 0.25 cm g
mass of dust). The quoted range of values correspond to variations in κ1mm = 0.25 from 0.5 to 0.1
2 −1
cm g .

38

a
Table 5
Molecular lines detected with SMAa
Species
HC3 N
CH3 OH
H2 CO
–
HC3 N v6 = 1
HC3 N v7 = 1
33 SO
2
OCS
HNCO
HC3 N v7 = 1
SO2
34 SO
2
CH3 CH2 CN
18 O
C
HNCO
–
–
H2 13 CO
SO
CH3 OH
CH3 OCHO
CH3 CHO vt = 1
13 CO
HNCO
CH3 CN
33 SO
2
CH3 CN
–
–
–
–
–
33 SO
2
34 SO
2
–
SO2
H
CH3 CCH
–
CH3 OCHO v = 1
CH3 CCH
H2 CCO
–

CO
OCS
13 CS
H
CH3 OCH3
CH3 OCH3
CH3 OH
–
CH3 CH2 CN
–
–
34 SO
2
CH3 OCHO
SO2

Transition

ν0 [GHz] Eu [K] Integrated flux [Jy km s−1 ]

LSB Low
24–23
218.3247
4(2,2)–3(1,2)
218.4400
3(2,2)–2(2,1)
218.4756
3(2,1)–2(2,0)
218.7600
24–23 l = 1f
218.8543
24–23 l = 1e
218.8608
22(2,20)–22(1,21)
218.8754
18–17
218.9033
10(1,10)–9(1,9)
218.9810
24–23 l = 1f
219.1737
22(7,15)–23(6,18)
219.2759
11(1,11)–10(0,10)
219.3550
22(2,21)–21(1,20)
219.4636
2–1
219.5603
10(2,9)–9(2,8)
219.7338
10(2,8)–9(2,7)
219.7371
10(0,10)–9(0,9)
219.7982
3(1,2)–2(1,1)
219.9084
6(5)–5(4)
219.9494
8(0,8)–7(1,6)
220.0784
17(4,13)–16(4,12) E
220.1668
10(3,8)–11(0,11) E
220.1800
LSB High
2–1
220.3986
10(1,9)–9(1,8)
220.5847
12(6)-11(6)
220.5944
11(1,11)–10(0,10)
220.6174
12(5)–11(5)
220.6410
12(4)-11(4)
220.6792
12(3)-11(3)
220.7090
12(2)-11(2)
220.7302
12(1)-11(1)
220.7430
12(0)-11(0)
220.7472
14(3,11)–14(2,12)
220.9857
22(2,20)–22(1,21)
221.1149
13(2,12)–13(1,13)
221.7357
11(1,11)–10(0,10)
221.9652
38β
222.0117
13(4)–12(4)
222.0991
13(3)–12(3)
222.1288
18(6,12)–17(6,11) A
222.1488
13(1)–12(1)
222.1627
11(0,11)–10(0,10)
222.1976
11(3,8)–10(3, 7)
222.2002
USB Low
2–1
230.5380
19–18
231.0609
5–4
231.2207
30α
231.9009
13(0,13)–12(1,12) AA
231.9877
13(0,13)–12(1,12) EA + AE 231.9879
USB High
10(2,8)–9(3,7) ++
232.4185
10(−3, 8) − 11(−2, 10)
232.9458
26(6,21)–25(6,20)
233.2050
26(6,20)–25(6,19)
233.2073
42(4,38)–41(5,37)
233.2910
10(5,5)–11(4,8)
233.2964
26(21,5)–27(20,7) E
233.7284
28(3,25)–28(2,26)
234.1870

Note

130.9
45.4
68.0
68.1
848.9
452.1
251.7
99.8
101.0
452.3
352.7
60.1
112.4
15.8
228.4
228.2
58.0
32.9
34.9
96.6
103.1
275.4

322
372
282
294
63
63
38
197
89
60
203
305
33
1.6943 × 104
80
80
198
63
3.042 × 103
118
93
63

15.86
101.5
325.8
61.0
247.3
183.1
133.1
97.4
76.0
68.8
120.2
248.1
92.5
60.3
–
190.2
139.6
312.3
81.7
63.9
181.3

1.23553 × 105
90
53
92
40
54
150
96
295
295
47
116
153
1397
159
33
118
106
427
13
13

(1)
(1)

16.5
110.8
33.3
–
80.9
80.9

–
255
202
747
122
122

(1)
(1)

165.4
190.3
190.9
190.9
410.1
109.8
499.2
403.0

90
60
47
47
33
33
75
591

(1)
(1)

(1)
(1)

(1)
(1)

(2)
(3)

(4)
(1)
(1)
(1)
(1)

a Molecular lines detected above an intensity I > 50 mJy beam−1 at line peak in the subcompact-array maps,
ν

averaged over an area equal to the 5σ contour of the continuum map shown in Figure 5, bottom. The first column
refers to the molecule tag, the second column to the transition, the third to its rest frequency as found in Splatalogue
(http://www.splatalogue.net/), the fourth to the upper-level energy, the fifth to the velocity-integrated flux in the
all-configuration maps, and the fifth to additional notes. Data used from Splatalogue are compiled from the CDMS
catalog (M¨ ller et al. 2005) and the NIST catalog (Lovas 2004). A single flux is listed for blended lines. The 13 CO
u
and C18 O 2–1 maps have been combined with IRAM30m single-dish data.
1 Blended lines.
2 From the measured line intensity of the H30α line, and assuming LTE conditions, this feature is dominated by
H38β. It could have a contribution of < 25% from CH3 CCH 13(6)–12(6), ν0 = 222.0144 GHz.
3 From the measured line intensity of the H30α line, and assuming LTE conditions and a 10% Helium abundance,
this feature is dominated by CH3 CCH 13(4)–12(4). It could have a contribution of < 20% from He38β, ν0 = 222.1022
GHz.
4 Could be contaminated with CH CHO 4(4,1)–5(3,2) A–, ν = 232.9518 GHz.
3
0

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