A multi-transition molecular line study of candidate massive young
  stellar objects associated with methanol masers
Astronomy & Astrophysics manuscript no. aa7289 c© ESO 2018
October 27, 2018
A multi-transition molecular line study of candidate massive
young stellar objects associated with methanol masers ⋆
M. Szymczak1, A. Bartkiewicz1, and A.M.S. Richards2
1 Toruń Centre for Astronomy, Nicolaus Copernicus University, Gagarina 11, 87-100 Toruń, Poland
2 Jodrell Bank Observatory, University of Manchester, Macclesfield, Cheshire SK11 9DL, UK
Received 13 February 2007 / Accepted 19 March 2007
ABSTRACT
Aims. We characterize the molecular environment of candidate massive young stellar objects (MYSOs) signposted by methanol
masers.
Methods. Single pixel observations of 10 transitions of HCO+, CO and CS isotopomers were carried out, using the IRAM 30 m
telescope. We studied a sample of 28 targets for which the 6.7 GHz maser emission positions are known with a sub-arcsecond accuracy.
Results. The systemic velocity inferred from the optically thin lines agrees within ±3 km s−1 with the central velocity of the maser
emission for most of the sources. About 64% of the sources show line wings in one or more transitions of CO, HCO+ and CS species,
indicating the presence of molecular outflows. Comparison of the widths of line wings and methanol maser emission suggests that the
6.7 GHz maser line traces the environment of MYSO of various kinematic regimes. Therefore conditions conducive for the methanol
maser can exist in the inner parts of molecular clouds or circumstellar discs as well as in the outer parts associated with molecular
outflows. Calculations of the physical conditions based on the CO and HCO+ lines and the CS line intensity ratios refine the input
parameters for maser models. Specifically, a gas number density of < 107 cm−3 is sufficient for strong maser emission and a high
methanol fractional abundance (> 5× 10−7) is required.
Key words. ISM: molecules − radio lines: ISM − stars: formation − masers
1. Introduction
There is compelling evidence that methanol masers are a
signature of recent or ongoing high-mass star formation
(Menten 1991). However, it is not yet fully understood when
they appear in an evolutionary sequence and what they actu-
ally trace. The evaporation of grain mantles is postulated as the
main process enhancing the fractional abundance of methanol
molecules in the gas phase up to 10−6 (Dartois et al. 1999). This
implies that methanol masers can emerge after the formation
of an embedded heating source. Methanol maser sources rarely
show strong (>100 mJy) free-free emission at centimetre wave-
lengths implying that they precede the development of detectable
ultra-compact HII (UCHII) region (Walsh et al. 1998; Codella &
Moscadelli 2000). The estimated lifetime of methanol masers is
a few ×104 yr (van der Walt 2005) which is similar to the typical
dynamical timescales of molecular outflows. High spatial reso-
lution observations revealed a variety of maser site sizes from
40−1200 AU (e.g. Norris et al. 1998; Walsh et al. 1998; Minier
et al. 2000). The maser emission arises either from circumstel-
lar discs or behind shocks tracing outflows from massive young
stellar objects (MYSOs). No object was found that unequivo-
cally confirms one of these scenarios.
The non-linear nature of maser amplification means that it is
difficult to relate the maser line intensity directly to the physical
parameters of the active region. Theoretical models predict the
formation of methanol maser lines under a rather wide range of
gas and dust temperatures (30−200 K and 100−300 K, respec-
tively) and hydrogen number densities (105− 108 cm−3) (Cragg
⋆ Figure A.1 and Table A.1 are only available in electronic form via
http://www.edpsciences.org
et al. 2002). Thus, it appears that a better understanding of the
environments in which the masers arise is required in order to
realise their full potential as probes of the formation of high-
mass stars.
In this paper we report our attempts to constrain the range
of environments probed by methanol masers using observations
of thermal emission from other molecular species and lines.
Specifically, the ratios of the intensities of different transitions
of CS and C34S molecules are used to obtain the temperature
and density of the gas. The optically thin and thick lines of CO
and HCO+ are used to constrain the column density. These tech-
niques were successfully used to characterize other samples of
MYSOs (e.g. Plume et al. 1997; Beuther et al. 2002a; Purcell et
al. 2006). Additionally, the molecular line profiles yield informa-
tion on the kinematics of various parts of the molecular clouds
surrounding the high-mass protostars (e.g. Fuller et al. 2005;
Purcell et al. 2006).
A homogeneous and unbiased sample of MYSOs is neces-
sary in order to address these issues properly. Our recent 6.7-
GHz unbiased survey for methanol masers in selected regions of
the Galactic plane (Szymczak et al. 2002) provides such a com-
plete, sensitivity limited sample of candidate MYSOs. Objects
identified in the survey probably represent a class of MYSOs in
an early evolutionary phase. Some groups and individual sources
in this class, selected using various diagnostics of high-mass star
formation, have been studied in thermal molecular lines (Brand
et al. 2001; Beuther et al. 2002a; Fuller et al. 2005), but this is
the first published study of a homogeneous sample based solely
on the presence of detectable methanol masers.
http://arxiv.org/abs/0704.1220v1
http://www.edpsciences.org
2 M. Szymczak et al.: Molecular line study of high-mass protostars
2. The sample
The 28 sources observed in this study (Table 1) were chosen
from a sample of 100 methanol maser sources found in the
Torun 32 m telescope blind survey for the 6.7 GHz methanol
line in the Galactic plane area 20◦ ≤ l ≤ 40◦ and |b| ≤ 0.◦52
(Szymczak et al. 2002). This flux-limited (3σ ≃1.6 Jy) subsam-
ple includes 25 out of 26 sources which were undetected prior to
the Torun survey. Therefore, our subsample specifically excludes
previously known sources associated with OH maser emission
(Caswell et al. 1995) or with IRAS-selected bright UCHII can-
didates (Schutte et al. 1993; van der Walt et al. 1995; Walsh et
al. 1997). Assuming that CH3OH masing precedes the appear-
ance of OH masers and detectable UCHII regions, the objects
studied here represent sites of high-mass star formation at a very
early stage. The average peak maser flux of the 28 targets is
17.3 Jy, a factor of 2.6 lower than that of the other 72 objects in
the original sample, suggesting that distant or intrinsically faint
objects may be over-represented in our subsample. The subsam-
ple studied here is most certainly not complete.
2.1. Astrometric positions
The coordinates and position uncertainties of the brightest 6.7-
GHz maser component in each source are presented in Table 1.
The LSR velocity of this component (Vp) and its peak flux den-
sity (Sp) are given for each target. The positions and flux den-
sities of all but three objects were measured with the Mark II −
Cambridge baseline of MERLIN in two sessions between 2002
May and 2003 May. For the three objects not measured the peak
velocities were taken from Szymczak et al. (2002).
The observational setup and data reduction were de-
scribed in Niezurawska et al. (2005). A primarily goal of
those astrometric measurements was to determine the positions
with sub-arcsecond accuracy for follow-up VLBI observations.
Measurement errors mainly depended upon the ratio of the beam
size to the signal to the noise ratio (Thompson et al. 1991). If the
emission was complex we took the dispersion of neighbouring
maxima as the position uncertainty. The N−S elongation of the
synthesized beam close to declination 0◦ produces a split peak,
in which case the position uncertainty in that direction was taken
as half the separation of the maxima. Consequently, for sources
with a single clear peak, the position errors in right ascension
were as small as 0.′′02 but increased up to 0.′′90 for sources with
complex emission. The respective errors in the declination were
0.′′14 and 1.′′6. Comparison with our unpublished VLBI observa-
tions reveals position differences between MERLIN and VLBI
measurements no worse than a few tens of milli-arcseconds. This
implies that the values listed in Table 1 are maximal position er-
rors for most of the targets. The flux densities listed in Table 1
are a factor of 2−3 lower than those measured with the single
dish (Szymczak et al. 2002) and should be considered as lower
limits. The exact flux scale and gain-elevation effects for low-
declination sources are not yet fully investigated at 6 GHz but
comparison of calibration sources in common with other exper-
iments shows that the uncertainties are 10 − ≤50%. This sug-
gests that about half the methanol flux arises on scales larger
than the beam size of 50−100 mas.
2.2. Distances
The distances were determined using the Galactic rotation curve
of Brand & Blitz (1993) and the central velocity of each 6.7 GHz
methanol maser profile as measured by Szymczak et al. (2002).
Selection of this velocity as a reliable estimator of the systemic
velocity is proven in Sect. 6.1. The sources are all in the first
quadrant so that there is an ambiguity between the near and far
kinematic distances. In most cases we are unable to resolve this
ambiguity because there are no independent distance measure-
ments in the literature for our sample. Based on the arguments
discussed in Walsh et al. (1997), we adopted the near kinematic
distances (Table 4).
3. Observations and data reduction
Observations were carried out between 2004 September 28 and
October 2 with the IRAM 30 m telescope. Ten transitions of
HCO+, CO and CS isotopomers were observed. Two or three
SIS receivers tuned to single sideband mode were used simulta-
neously, in combination with the VESPA autocorrelator as well
as with 100 kHz and 1 MHz filter banks. Table 2 lists the rest line
frequencies, half power beam widths (HPBWs), velocity resolu-
tions and typical system temperatures for each transition.
The data were taken using the position switching mode. The
off positions were typically 30′ away from the targets. In the
few cases, especially for the C13O J=2−1 line, where emission
was seen at the reference position, the offsets were increased up
to 45′ in the direction away from the Galactic plane. The ob-
servations were centered on the target positions listed in Table
1. Integration times on-source in total power mode were 12−18
min per frequency setting, resulting in rms noise levels ranging
from ≈0.05 K at 87 GHz to ≈0.90 K at 245 GHz for a spectral
resolution of 0.10−0.16 km s−1. Pointing was checked regularly
on nearby continuum sources and was usually found to be within
2′′ and always within 3′′. The spectra were scaled to the main
beam brightness temperature (Tmb) using the efficiencies sup-
plied by the observatory1. Comparison of our data with those
taken by Brand et al. (2001) with the same telescope and spec-
tral resolutions for a source in common, 36.115+0.552, implies
consistent intensity scales within less than 30%.
The data reduction were performed using the CLASS soft-
ware package. Low order polynomials were applied to remove
baselines from the calibrated spectra. The line parameters were
determined from Gaussian fits and are listed in Table A.1 where
the following information is given: the rms (1σ) noise level, the
extreme velocities Vs, Ve where the intensity drops below the 2σ
level, the peak temperature Tmb, the velocity of the peak Vp, the
line width at half maximum ∆V and the integrated line intensity∫
Tmbdv. Velocities are in the LSR frame. In several cases where
the profiles were non-Gaussian, these values were read off from
the spectra. In some cases the spectra were smoothed to improve
the signal to noise ratio. In this paper, only the autocorrelation
spectra are analyzed.
4. Results
The basic parameters of molecular transitions derived from
Gaussian analysis are assembled in Table A.1, while the spec-
tra are shown in Fig. A.1.
The sensitivity achieved and detection rate for each transition
are summarized in Fig. 1. The histogram counts as detected only
those lines with Tmb > 3σ.
13CO(2−1), C18O(2−1), CS(2−1)
and CS(3−2) lines were detected in all sources. HCO+(1−0)
and H13CO+(1−0) lines were detected in all but one source. The
detection rates in C34S(2−1) and C34S(3−2) transitions were
1 http://www.iram.es/IRAMES/telescope/telescopeSummary/telesco-
pe summary.html
M. Szymczak et al.: Molecular line study of high-mass protostars 3
Table 1. List of targets.
Name α(J2000) δ(J2000) σα σδ Vp Sp
(arcsec) (arcsec) (km s−1) (Jy)
21.407−0.254 18 31 06.3403 −10 21 37.305 0.28 0.80 +89.0 2.0
22.335−0.155 18 32 29.4109 −09 29 29.435 0.27 1.10 +35.7 2.8
22.357+0.0661 18 31 44.144 −09 22 12.45 +80.
23.707−0.1983 18 35 12.3625 −08 17 39.409 0.06 0.40 +79.0 3.2
23.966−0.1093 18 35 22.2167 −08 01 22.395 0.35 1.60 +71.0 4.3
24.147−0.0093 18 35 20.9501 −07 48 57.470 0.03 0.19 +17.9 6.4
24.541+0.3123 18 34 55.7212 −07 19 06.630 0.90 0.90 +105.5 4.4
24.635−0.323 18 37 22.7932 −07 31 37.950 0.50 1.20 +35.5 1.0
25.410+0.1052 18 37 16.9 −06 38 30.4 +97.
26.598−0.024 18 39 55.9268 −05 38 44.490 0.03 0.18 +23.0 2.0
27.221+0.136 18 40 30.5446 −05 01 05.450 0.03 0.18 +119.0 3.0
28.817+0.365 18 42 37.3470 −03 29 41.100 0.02 0.18 +91.0 1.0
30.316+0.069 18 46 25.0411 −02 17 45.160 0.03 0.16 +35.5 1.3
30.398−0.297 18 47 52.2623 −02 23 23.660 0.02 0.14 +98.2 1.5
31.056+0.361 18 46 43.8558 −01 30 15.690 0.05 0.28 +81.0 1.0
31.156+0.045 18 48 02.3471 −01 33 35.095 0.10 0.90 +41.0 0.8
31.585+0.080 18 48 41.8975 −01 09 43.085 0.50 0.70 +95.8 0.8
32.966+0.0412 18 51 24.5 +00 04 33.7 +92.
33.648−0.2243 18 53 32.5508 +00 32 06.525 0.50 1.0 +62.6 20.0
33.980−0.019 18 53 25.0184 +00 55 27.260 0.05 0.50 +59.0 1.0
34.753−0.092 18 55 05.2410 +01 34 44.315 0.08 0.50 +53.0 1.6
35.791−0.1753 18 57 16.9108 +02 27 52.900 0.04 0.17 +60.8 5.6
36.115+0.5523 18 55 16.8144 +03 05 03.720 0.02 0.23 +74.2 7.2
36.704+0.096 18 57 59.1149 +03 24 01.395 0.08 0.17 +53.0 1.9
37.030−0.039 18 59 03.6435 +03 37 45.140 0.14 0.50 +79.0 1.2
37.479−0.105 19 00 07.1457 +03 59 53.245 0.07 0.36 +62.8 1.8
37.600+0.426 18 58 26.8225 +04 20 51.770 0.03 0.70 +91.2 2.0
39.100+0.4913 19 00 58.0394 +05 42 43.860 0.34 0.17 +15.2 2.8
1 Position is from Walsh et al. 1998, 2 Position is from Beuther et al. 2002a, 3 This source was reported in Niezurawska et al. 2005
Table 2. Observing parameters
Transition Frequency Ref. HPBW Res. Tsys
(GHz) (′′) (km s−1) (K)
HCO+(1−0) 89.188518 2 27 0.13 200
H13CO+(1−0) 86.754330 1 27 0.13 200
13CO(2−1) 220.398686 2 11 0.11 1200
C18O(2−1) 219.560328 2 11 0.11 1340
CS(2−1) 97.980953 1 25 0.12 260
CS(3−2) 146.969049 2 17 0.16 690
CS(5−4) 244.935560 1 10 0.10 1830
C34S(2−1) 96.412961 1 25 0.12 290
C34S(3−2) 144.617147 2 17 0.16 560
C34S(5−4) 241.016113 1 10 0.10 2100
The references for the line frequencies are 1 - Lovas (2003), 2 - Brand
et al. (2001)
about half of those in CS(2−1) and CS(3−2) lines. Because the
sensitivities achieved for these four lines were comparable, these
detection rate differences reflect a real drop in the number of
sources exhibiting emission at the same level in the C34S(2−1)
and C34S(3−2) lines. In contrast, the lower detection rates in the
CS(5−4) and C34S(5−4) transitions appear to reflect the drop in
sensitivity.
4.1. Systemic velocities
Five of the observed lines (C18O(2−1), H13CO+(1−0) and the
J=2−1, 3−2 and 5−4 transitions of C34S) are expected to be
optically thin (Plume et al. 1997; Brand et al. 2001; Purcell et
1/2520/2716/2819/1914/2525/2517/1827/2825/25
C34S 
(5-4)
(5-4)
C34S 
(3-2)
(3-2)
13CO 
(2-1)
C18O 
(2-1)
HCO+ 
(1-0)
H13CO+
 (1-0)
(2-1)
 C34S 
(2-1)
28/28
Fig. 1. The average sensitivity achieved for each transition (top)
and the detection rate (bottom). The ratio of the number of de-
tected to observed objects is shown in each of the bars.
al. 2006). These lines can be used to determine source systemic
velocities. In order to test whether these species trace the same
or similar kinematic regimes we compare their line parameters.
The C34S(5−4) line is excluded from the following analysis due
to very low number of detections.
The average intensities of the H13CO+(1−0) and C34S lines
are very similar and are a factor of 8 weaker than the average
4 M. Szymczak et al.: Molecular line study of high-mass protostars
intensity of the C18O(2−1) line. This implies that the parameters
of the latter line, especially Vp, are determined most accurately.
We note that some line rest frequencies adopted from Brand
et al. (2001) differ slightly from those recommended by Lovas
(2003). In the extreme case of C34S(3−2) this results in the ve-
locity difference of 0.07 km s−1. Moreover, the uncertainties in
the line rest frequencies introduce a maximum uncertainty of
±0.17 km s−1 to the derived line velocity for the H13CO+(1−0).
We assume that the above uncertainties affect the velocity esti-
mates by up to 0.24 km s−1. Comparison of the velocities of the
four optically thin lines in our sample reveals no significant aver-
age differences higher than 0.30 km s−1. This suggests the same
kinematic behaviour of these low-density gas tracers.
At 100 K the thermal linewidths of C18O(2−1) and
H13CO+(1−0) are 0.24 km s−1 whereas those of C34S(2−1)
and C34S(3−2) are 0.20 km s−1. The observed linewidths are
much broader, suggesting that turbulence or bulk gas motions
play a significant role in the line broadening. The mean linewidth
ratios of the optically thin lines are 5−10% higher than unity.
This bias is relatively small and suggests that the lines trace
the same molecular gas in the beam. The systemic velocities
are listed in Table 4. They are primarily the C34S(2−1) and
C34S(3−2) line peak velocities. If emission in these lines is
absent or weak the other optically thin lines are used. In two
sources, 37.030−0.039 and 37.600+0.426, the systemic veloc-
ities are derived from CS(5−4) and HCO+(1−0) profiles, re-
spectively. We conclude that in most cases the observed opti-
cally thin lines are well fitted by single Gaussian profiles (devi-
ations are discussed in Sect. 4.2.2) and their peak velocities are
within ±0.4 km s−1 of each other for almost all sources in the
sample. Therefore, these lines provide reliable estimates of sys-
temic velocity of sufficient accuracy to allow comparison with
the methanol maser velocities listed in Szymczak et al. (2002).
4.2. Shape of profiles
We analyse the molecular line profiles in order to search for
specific signatures of ordered motions such as infall, outflow or
rotation. Inward motion can be signposted by blue asymmetric
profiles (Myers et al. 1996; Fuller et al. 2005) if the molecular
spectral lines trace sufficiently dense gas. Signatures of outflow
or rotation are generally manifested in the line wings.
4.2.1. Asymmetry
We analysed line asymmetry quantitatively using the asymmetry
parameter (Mardones et al. 1997), δv=(vthick − vthin)/∆Vthin,
where vthick and vthin are the peak velocities of optically thick
and optically thin lines, respectively and ∆Vthin is the line
width at half maximum of the optically thin line. We used
C34S(2−1) as the optically thin line and the best available mea-
sure of the systemic velocity of MYSOs. Figure 2 shows his-
tograms of the distribution of δv for the optically thick lines
13CO(2−1), HCO+(1−0), CS(2−1), CS(3−2) and CS(5−4).
There are approximately equal numbers of blue and red asym-
metric profiles in our sample. Specifically, we note that there is
no evidence for an excess of blue-shifted emission in the opti-
cally thick lines. Such an excess is postulated as the signature
of inward motion of the gas (Myers et al. 1996). We suggest
that motions other than infall, i.e. turbulence, rotation and out-
flow, are the dominant factor producing asymmetric profiles for
most of the sources in our sample. It is possible that infall sig-
natures could be masked by the relatively low resolution (typi-
-1 0 1
-1 0 1 -1 0 1 -1 0 1 -1 0 1
13CO(2-1)
HCO+(1-0)
CS(2-1)
CS(3-2)
CS(5-4)
Fig. 2. Histograms of the distribution of the asymmetry parame-
ter δv for the five transitions. The range of |δv| < 0.25 marked
by the dotted lines corresponds to the spectra with no asymme-
cally ≥ 0.2 pc, i.e. at a distance of 5 kpc and spatial resolution
of 10′′) of our observations, since even the near kinematic dis-
tances are >3 kpc for ∼80% of the sources (the average Dnear
is 5.2± 2.5 kpc for the whole sample).
We therefore examined separately the 5 closest (Dnear <
2.8 kpc) objects with well-determined asymmetry parameters.
Two of these, 26.598−0.024 and 30.316+0.069, consistently
show negative values of δv, i.e. blue asymmetry, in the
13CO(2−1), HCO+(1−0), CS(2−1) and CS(3−2) line pro-
files (Fig. A.1). The corresponding values of δv are −0.38,
−0.82, −0.15 and −0.18 for source 26.598−0.024 and −0.43,
−0.60, −0.33 and −0.37 for source 30.316+0.069. Their asym-
metry parameters are smaller for the optically thin lines (i.e.
C18O(2−1) and H13CO+(1−0)), in the range from −0.20 to 0.0.
Such a dependence of the amount of blue asymmetry on the op-
tical depth of the transition is typical in molecular cores experi-
encing infall (Narayanan et al. 1998). We suggest that these two
sources are the clearest infall candidates although source com-
plexity or a combination of outflow and rotation could contribute
to asymmetries in optically thick profiles.
4.2.2. Wings
Wing emission is identified by the presence of residuals after
Gaussian fitting and by comparing the same transitions of op-
tically thick and thin isopotomers. A single Gaussian function
provides a good fit to most of the optically thin lines analyzed in
Section 4.1, but in a few cases the residuals are at a level ≥ 3σ,
assumed to be wing emission. We cannot exclude the possibility
that they are weak separate component(s), given the limitations
of our signal to noise ratio and spectral resolution, but we note
that the blue and/or red residuals are non-Gaussian in most cases.
4 out of 25 sources detected in the C18O(2−1) line show weak
(3-4σ) wing emission of width 4.5−8 km s−1 (Table 3) which
mostly is seen from the red or blue sides of the profiles. In the
H13CO+(1−0) line the wing emission is seen in 2 out of 17 ob-
jects detected (Table 3). 25.410+0.105 is a peculiar source show-
ing broad (7−8 km s−1) and symmetric wings in both lines (Fig.
A.1).
In contrast, the optically thick lines show more frequent ab-
sorption dips, multiple components and wings. In several cases
identification of the wings is difficult. The 13CO(2−1) profiles
are especially complex; commonly they are fit by 2-5 Gaussians.
These profiles can be interpreted as multiple emitting regions
M. Szymczak et al.: Molecular line study of high-mass protostars 5
Table 3. Statistics of wing occurrence. Entries marked Y or N
indicate symmetric wings and no wings respectively, Yb or Yr
indicate wing emission seen from the blue and red sides of the
profiles, respectively. An interrogation point (?) indicates a ten-
tative wing and the absence of entry indicates no observation.
Source C18O H13CO+ 13CO HCO+ CS CS
(2-1) (1-0) (2-1) (1-0) (2-1) (3-2)
21.407−0.254 N N N
22.335−0.155 N N N Y N N
22.357+0.066 Yb N N N N N
23.707−0.198 N N N N N N
23.966−0.109 N Yr N Y Y Y
24.147−0.009 N N N Y N N
24.541+0.312 Y N N Y N N
24.635−0.323 Yr N Y Y Y Y
25.410+0.105 N Y Y Y Y Y
26.598−0.024 N N N N N N
27.221+0.136 N N N N
28.817+0.365 N N Y Y
30.316+0.069 N N N Y N N
30.398−0.297 Yb N ? Y Y N
31.056+0.361 N N N N
31.156+0.045 N N ? Y Y Y
31.585+0.080 N N N N
32.966+0.041 N N ? N N N
33.648−0.224 N N N N N
33.980−0.019 N N N Y Y Y
34.753−0.092 N ? Y N
35.791−0.175 N N ? Y Y Y
36.115+0.552 N N N Y N N
36.704+0.096 N N Y N N
37.030−0.039 N N
37.479−0.105 ? N
37.600+0.426 Y Y
39.100+0.491 N N N Y Y Y
along the same line of sight. The 13CO(2−1) lines show evi-
dence of wings in only 3 objects (Table 3).
The HCO+(1−0) lines are also complex, often exhibiting
two or more components or broad line wings (Fig. A.1). They
appear to consist of the superposition of several emitters seen
along the line of sight or of (self)absorption by cooler gas on
the near side of the source. Wings are identified in 17 out of
27 detections (Table 4). The wing full width ranges from 6 to
20 km s−1 with a mean value of 10.3±3.3 km s−1.
Evidence for wings is seen in the CS(2−1) transition for 9
out of 25 sources and in the CS(3−2) transition for 7 out of 19
sources (Table 3). Their full widths are from 8 to 19 km s−1.
We conclude that 64% (18/28) of the sources show residual
line wings at least in one line when a Gaussian profile is used
to fit the CO, HCO+ and CS molecular lines. Detection of the
wings may indicate molecular outflows from the MYSOs iden-
tified by methanol masers but we caution that such detections
based on our data alone are only tentative.
5. Derivation of physical parameters
5.1. Column densities
In order to estimate the column density of H13CO+ from the ob-
served HCO+(1−0) and H13CO+(1−0) line parameters, we fol-
low the procedure outlined in Purcell et al. (2006) and references
therein. Briefly, the main assumptions made are: (i) HCO+(1−0)
is optically thick and H13CO+(1−0) is optically thin. (ii) Both
lines form in the same gas and share the same excitation temper-
ature. (iii) The excitation temperature is equal to the rotational
temperature. (iv) The gas is in local thermodynamic equilibrium.
(v) The beam filling factor is one for both lines.
The derived H13CO+ column density, N (H13CO+), (Table
4) ranges from 1.3 − 5.1 × 1012 cm−2 and the median value
is 2.2 × 1012 cm−2. We derive a value of N (H13CO+) a fac-
tor of 4 smaller than the value found by Purcell et al. (2006)
for two of the sources common to both samples, 22.357+0.066
and 23.707−0.198. This is probably because Purcell et al. ap-
plied corrections for self-absorption, leading to higher estimates
of the HCO+(1−0) line intensities and lower optical depths,
compared with our study. We adopt an abundance ratio of
[H13CO+/H2]=3×10
−11 (Girart et al. 2000), from which we ob-
tain the H2 column density from 4.3 − 17.0 × 10
22 cm−2 with
the median value of 7.3× 1022 cm−2.
We apply the same method to estimate the column density
of C18O, N (C18O), from the line parameters of 13CO(2−1)
and C18O(2−1), assuming that 13CO(2−1) is optically thick
and C18O(2−1) is optically thin. For our sample N (C18O)
is 0.9−32.6×1015 cm−2 (Table 4) with the median value of
4.6×1015 cm−2. The temperature varies between 10 and 30 K.
The resulting H2 column density ranges from 5.4 × 10
1.9×1023 cm−2 for an abundance ratio [C18O]/[H2]=1.7×10
(Frerking et al. 1982).
We conclude that the CO and HCO+ data provide consis-
tent estimates of the column density of H2 towards the methanol
maser sources. The range of N (H2) derived here is in good
agreement with that reported for high-mass protostar candidates
associated with methanol masers; 3 × 1022 − 2 × 1023 cm−2
(Codella et al. 2004; Minier et al. 2005; Purcell et al. 2006).
However, it is significantly lower than N (H2)≥ 4 × 10
23 cm−2
reported in some earlier works (e.g. Churchwell et al. 1992) for
ultra-compact HII regions. This discrepancy is likely due to the
temperatures of 10−30 K derived here which is significantly
lower than ≥ 100K assumed in Churchwell et al. (1992).
We notice that a dispersion of the N (C18O) is a factor of 7
larger than that of the N (H13CO+) (Table 4). In two sources
22.357+0.066 and 26.598−0.024 the N (C18O) is extremely
large (> 1.9×1016 cm−2). In consequence, the values of N (H2)
derived from the C18O is a factor of 1.5 and 3.1, respectively,
higher that those derived from the H13CO+. This discrepancy
suggests that the methanol masers in these sources probe re-
gions with the abundance ratio of 13CO/C18O significantly lower
than a typical ratio of 6.5−7 (Frerking et al. 1982; Beuther et
al. 2000). A decrease of 13CO/C18O ratio is predicted in the PDR
model in a clumpy cloud; in small clumps the C18O molecule is
nearly completely photodissociated whilst it is protected from
photodissociation in large clumps (Beuther et al. 2000 and ref-
erences therein). Object 26.598−0.024 with the highest value of
N (C18O) is also a candidate infall object (Sect. 4.2.1) and one
can speculate that it is the youngest methanol maser in our sam-
ple; the maser emission forms in large clumps at nearly systemic
velocity. Another explanation of low 13CO/C18O intensity ratio
can be that our 11′′ beam probes the methanol maser sites where
the C18O cores did not coincide with the13CO cores. This obser-
vational fact is well documented in Brand et al. (2001) at least for
their sources Mol 98 and Mol 136 (see their Fig. 5). Furthermore,
the C18O emission is less extended than the 13CO emission; by
a factor of ∼ 3 − 5 for common source 35.791−0.175. This ex-
planation seems to be less plausible as a similar effect can be
observed for HCO+ and H13CO+ lines.
6 M. Szymczak et al.: Molecular line study of high-mass protostars
Table 4. Derived properties
30K 60K
Source Vsys dnear dfar N (H
13CO+) N (C18O) lognH2 logN (CS) lognH2 logN (CS)
(km s−1) (kpc) (kpc) (1012cm−2) (1015cm−2) (cm−3) (cm−2) (cm−3) (cm−2)
21.407−0.254 90.7 6.0 10.4 - 3.8 - - - -
22.335−0.155 30.9 2.4 14.7 2.1 3.7 6.15±0.15 14.52±0.10 5.91±0.12 14.68±0.06
22.357+0.066 84.2 5.2 10.6 2.2 19.1 5.48±0.09 14.70±0.11 5.27±0.13 14.52±0.21
23.707−0.198 68.9 5.1 10.5 3.2 13.2 5.93±0.08 13.74±0.14 5.57±0.08 13.49±0.07
23.966−0.109 72.7 4.2 11.6 5.1 9.0 >6.7 14.73±0.57 >6.5 15.16±0.52
24.147−0.009 23.1 2.0 14.5 1.4 2.1 5.61±0.08 14.51±0.18 5.42±0.12 14.60±0.56
24.541+0.312 107.8 7.0 9.5 1.5 4.6 - - - -
24.635−0.323 42.7 3.7 13.1 4.7 7.6 >6.7 14.61±0.38 6.39±0.12 14.53±0.23
25.410+0.105 96.0 - 9.5 3.4 7.0 6.42±0.11 14.40±0.08 6.22±0.07 14.53±0.20
26.598−0.024 23.3 1.8 13.4 1.8 32.6 >6.9 14.54±0.42 >6.5 14.86±0.28
27.221+0.136 112.6 - 8.0 - 9.4 - - - -
28.817+0.365 87.0 5.5 9.4 - 5.3 - - - -
30.316+0.069 45.3 2.8 12.2 2.2 3.3 >6.9 14.77±0.18 6.28±0.19 14.59±0.09
30.398−0.297 102.4 6.0 8.5 1.6 3.7 6.12±0.10 14.83±0.08 - -
31.056+0.361 77.6 - 9.6 - 2.9 - - - -
31.156+0.045 38.9 2.7 11.9 2.2 4.8 6.06±0.04 14.11±0.06 5.74±0.06 14.64±0.13
31.585+0.080 96.0 5.4 8.1 - 11.8 - - - -
32.966+0.041 83.4 5.4 8.9 1.3 4.2 - - 4.39±0.11 15.73±0.16
33.648−0.224 61.5 - 10.4 - 2.1 - - - -
33.980−0.019 61.1 3.5 10.6 2.5 4.7 - - 4.52±0.21 15.76±0.13
34.753−0.092 51.1 3.1 11.0 - 1.4 - - - -
35.791−0.175 61.9 4.6 10.3 2.4 3.0 - - - -
36.115+0.552 76.0 4.9 9.0 1.9 8.1 - - >6.890 15.60±0.1990
36.704+0.096 59.8 4.6 10.4 - 0.9 - - - -
37.030−0.039 80.1 5.0 8.3 - - - - - -
37.479−0.105 59.1 - 9.5 - - - - - -
37.600+0.426 90.0 6.5 7.5 - - - - - -
39.100+0.491 23.1 1.0 14.7 2.0 2.9 - - 6.58±0.0890 14.81±0.0890
90 values for kinetic temperature 90 K
5.2. Gas density and temperature
We used the escape-probability modelling code RADEX on-
line2 to estimate the density and temperature of the gas required
for the observed line temperature ratios of CS and C34S. Because
these parameters cannot be derived independently for diatomic
molecules (Schilke et al. 2001) we calculate the models for 30,
60 and 90 K with gas number densities of 104 − 107 cm−3, CS
column densities of 1012−1017 cm−2 and linewidth of 1 km s−1.
We performed the calculations for the 16 sources for which all
three CS lines were detected and we assumed that beam dilu-
tion is comparable for all these transitions. We used a χ2 mini-
mization procedure to fit the models to the observed line ratios.
The derived parameters are listed in Table 4. We found equally
reasonable fits for 10 sources using models at kinetic tempera-
tures of both 30 and 60 K. Five sources have good fits only for
a single kinetic temperature. We could not find a satisfactory fit
for the source 35.791−0.175 as its CS(2−1) and CS(3−2) lines
are strongly self-absorbed (Fig. A.1) and thus its line ratios are
poorly constrained.
Using a temperature of 60 K the average logarithmic num-
ber density is 5.7±0.7 and the average logarithmic column den-
sity of CS is 14.7±0.6 for the sample. These values are consis-
tent with 5.9 and 14.4, respectively, reported for a large sam-
ple of massive star formation sites selected by the presence of
H2O masers (Plume et al. 1997). Our estimates are also in good
agreement with those based on the nine-point CS maps of high-
mass protostellar candidates (Beuther et al. 2002a; Ossenkopf et
al. 2001) and calculated with more sophisticated models. Taking
2 http://www.strw.leidenuniv.nl/moldata/radex.php
the CS fractional abundance as ∼ 8×10−9 (Beuther et al. 2002a)
our estimate of the CS column density implies a mean N (H2) of
6.3× 1022 cm−2 which is in very good agreement with the esti-
mates based on CO and HCO+ data (Sect. 5.1).
Our C34S data are less useful to estimate the gas density and
temperature because the line ratios are poorly constrained for
most of the targets. 26.598−0.024 is the only source for which
we are able to determine C34S line ratios but the results are in-
consistent with those obtained from the CS data. This indicates
that the escape probability model provides only a crude esti-
mate to the physical parameters and the assumption of homo-
geneous parameters across the cloud is not fulfilled (Ossenkopf
et al. 2001).
6. Discussion
6.1. Kinematics
The present survey reveals new information regarding the kine-
matics of molecular gas surrounding massive forming stars. In
the following we attempt to answer the question of whether the
6.7 GHz methanol maser and the thermal molecular lines arise
from similar or different kinematic regimes.
The velocity ranges of 6.7 GHz methanol masers, 13CO
and HCO+ line wings are plotted in Fig. 3. This plot clearly
shows that the systemic velocity derived in this study (Table 4)
is in good agreement with the methanol maser central veloci-
ties, Vm, derived from Szymczak et al. (2002). We note that in
many sources Vm does not coincide with the peak maser veloc-
ity Vp. The average value of Vm − Vsys is 0.04±0.60 km s
The difference is less than 3 km s−1 for 23 sources (82%). Vm
M. Szymczak et al.: Molecular line study of high-mass protostars 7
-20 -10 0 10 20
39.100+0.491
37.600+0.426
37.479−0.105
37.030−0.039
36.704+0.096
36.115+0.552
35.791−0.175
34.753−0.092
33.980−0.019
33.648−0.224
32.966+0.041
31.585+0.080
31.156+0.045
31.056+0.361
30.398−0.297
30.316+0.069
28.817+0.365
27.221+0.136
26.598−0.024
25.410+0.105
24.635−0.323
24.541+0.312
24.147−0.009
23.966−0.109
23.707−0.198
22.357+0.066
22.335−0.155
21.407−0.254
 Velocity (km s-1)
Fig. 3. Comparison between the velocity ranges of 6.7 GHz
methanol maser (thick bars)(Szymczak et al. 2002) and 13CO
(dotted bars) and HCO+ (dashed bars) line wings. The dotted
vertical line marks the systemic velocity.
is offset by >4 and ≤8.1 km s−1 with respect to Vsys in 5
sources (18%), 23.707−0.198, 23.966−0.109, 24.147−0.009,
30.316+0.069 and 32.966+0.041 (Figs. 3 and A.1). This does
not necessarily imply that the different species arise from sep-
arate regions along the same line of sight. Two of the sources,
24.147−0.009 and 32.966+0.041, have ranges of maser emis-
sion ∆Vm ≤4 km s
−1 which is a factor of two narrower than
the mean value of 8.3±0.9 km s−1 for the sample but this could
be simply an effect of inhomogeneous conditions in molecular
clumps; the maser emission is sustained in one or a few clumps
of sizes a few×1015 cm (Minier et al. 2000). The effect of
clumping is clearly seen even in regular structures (Bartkiewicz
et al. 2005). The other three sources exhibit maser emission at
velocities which differ from the systemic velocity by less than
4 km s−1. In source 30.316+0.069 the maser spectrum is double
(Szymczak et al. 2000) and one of the peaks near 49 km s−1 is
close to the systemic velocity of 45.3 km s−1, so that the maser
emission related to the thermal molecular lines has a width of
about 6 km s−1. We conclude, Vm is a reliable estimator of the
systemic velocity, with an accuracy better than 3 km s−1, for
most of the sources in our sample.
The overlap between the velocity ranges of the methanol
masers and the 13CO/HCO+ line wings is remarkable. Figure
4 shows a histogram of the ratio of methanol maser velocity
spread, ∆Vm, to HCO
+ line wings spread. This ratio ranges
from 0.2−6.7 and the median value is 1.3. Similar trends are ob-
served in the ratio of ∆Vm to
13CO line wings spread. In 12 out
of 23 sources where we detected 13CO/HCO+ wings, ∆Vm falls
entirely within the wing velocity ranges and in 9 sources there is
an overshoot of ≤4 km s−1. The 13CO/HCO+ line wings appear
to provide a good indication of the presence of outflow and their
widths can serve as an approximate measure of outflow veloci-
ties. The present observations used beamwidths of 11′′ and 27′′
for 13CO and HCO+ lines, respectively, which samples a small
fraction of the molecular cloud, centred on the methanol maser
position.
The outflow velocity can be reliably estimated from these
data only for the fortunate case when the axis of out-
flow lies along the line of sight. One source in our sam-
ple, 25.410+0.105, has been mapped in the 12CO(2−1) line
by Beuther et al. (2002b) who measured a wing velocity
range of 14 km s−1, which is comparable with our estimate.
In this object the maser emission, with velocity width of
5 km s−1, is closely centered on the systemic velocity. The
velocity ranges of the 13CO and HCO+ wings are 11 and
18 km s−1, respectively. This indicates that the maser emission
traces a small portion of the kinematic regime of the 13CO
and HCO+ lines or it is completely unrelated. Fig. 3 indi-
cates that sources 21.407−0.254, 26.598−0.024, 31.156+0.045
and 35.791−0.175 share similar properties with 25.410+0.105.
VLBI observations of 35.791−0.175 (Bartkiewicz et al. 2004)
support the above interpretation. In this object the 6.7 GHz
methanol maser emission appears to come from part of a cir-
cumstellar disc.
Our sample contains 4 objects (23.707−0.198,
24.147−0.009, 32.966+0.041, 36.115+0.552) for which
the velocity range of the maser emission is very similar to
or slightly overshoots that of the 13CO/HCO+ line wings. If
we assume that the width of 13CO and HCO+ line wings is a
measure of the outflow velocity, in these objects the 6.7 GHz
methanol masers arise in outflows. This scenario appears to be
supported by VLBI observations of 36.115+0.552 (Bartkiewicz
et al. 2004); the maser emission comes from two well separated
regions which probably represent a bipolar outflow. In this case
the methanol maser traces the same or a very similar kinematic
regime as that of the 13CO and HCO+ lines.
Sources 22.355−0.155 and 27.221+0.136 appear to posses
complex kinematics in the regions where the methanol masers
operate. A close inspection of their 6.7 GHz spectra (Szymczak
et al. 2002) suggests that some spectral features arise from the
inner parts of the molecular cloud whilst other features form in
outflows. VLBI studies of maser emission and detailed measure-
ments of the kinematic properties of the molecular emission are
needed to verify this suggestion.
6.2. Implications for the evolutionary status
One of the important findings of our observations is the de-
tection of considerable number of sources with line wings. We
identified residual line wings in 18 out of 28 sources when a
Gaussian profile was used to fit the CO, HCO+ and CS molecu-
lar lines. The line wings appear to be the best indicators of out-
flow motions in most cases. The presence of line wings in about
64% of sources in the sample suggests a close association of the
methanol masers with the evolutionary phase when outflows oc-
cur. This result is consistent with that reported by Zhang et al.
(2005). They mapped the CO(2−1) line in a sample of 69 lu-
minous IRAS point sources and found that about 60% of them
were associated with outflows. However, with the present data
8 M. Szymczak et al.: Molecular line study of high-mass protostars
0 1 2 3 4 5 6 7 8
Ratio
Fig. 4. Histogram of the ratio of methanol maser velocity spread
to HCO+ line wings spread.
we cannot resolve whether the methanol maser sites and the out-
flows have a common origin. Because of clustering in high mass
star formation (e.g. Beuther et al. 2002a) it is possible that some
masers in the sample are not actually associated with outflowing
sources.
Codella et al. (2004) proposed an evolutionary sequence for
UCHII regions in which the earliest phase is marked by maser
emission and molecular outflows not yet large enough to be de-
tected with single-dish observations. The present data suggest
that our sources are slightly more evolved because several of
them show evidence of outflows. Their age therefore seems to
be less than a few 104 yr (Codella et al. 2004) which is consis-
tent with a statistical estimate of 3− 5× 104 yr for the methanol
maser lifetime (van der Walt 2005).
6.3. Constraints on maser models
The present study allows us to refine the range of physical con-
ditions required to produce strong methanol masers at 6.7 GHz.
Theoretical modelling by Cragg et al. (2002) demonstrated that
a maser line of 1 km s−1 width attains a peak brightness temper-
ature of ∼1011 K for a dust temperature >100 K and a methanol
column density > 5 × 1015 cm−2. They found that methanol
masers can be produced under a wide range of the physical
conditions. In fact, for a methanol fractional abundance from
3× 10−8 to 10−5, masing is predicted for the gas density range
5 − 2 × 108 cm−3 and the methanol column density range
5 × 1015 − 2 × 1018 cm−2 (Cragg et al. 2002). The gas density
inferred from our observations is between 105 and 107 cm−3;
higher values (> 107 cm−3) are less probable. The hydrogen
column density from 1022 to 2×1023 cm−2, inferred here, trans-
lates well into the above range of methanol column densities for
methanol fractional abundances of 5 × 10−7 − 10−5. This sug-
gests that 6.7 GHz maser emission is less probable in environ-
ments with a lower methanol fractional abundance of the order
of 10−8. We conclude that our study well refines a range of the
input parameters of Cragg et al.’s maser model. Specifically, a
high methanol fractional abundance of > 5 × 10−7 is required
whilst a gas density < 107 cm−3 is sufficient for the production
of methanol masers.
7. Conclusions
We have observed 10 transitions of HCO+, CO and CS iso-
topomers at millimetre wavelengths in order to characterize the
physical conditions in a sample of 28 MYSOs identified by the
presence of methanol masers. No other preconditions were in-
volved in the sample selection. The observations were centred at
maser positions known with a sub-arcsecond accuracy. The main
conclusions of the paper are summarized as follows:
(1) The systemic velocity determined from the optically
thin lines C18O(2−1), H13CO+(1−0), C34(2−1) and C34(3−2)
agrees within ±3 km s−1 with the central velocity of the
methanol maser emission for almost all the sources.
(2) 18 out of 28 sources show residual line wings at least
in one line when a Gaussian function was used to fit the CO,
HCO+ and CS lines. Detection of the line wing emission sug-
gests the presence of molecular outflows in these sources. Their
occurrence needs to be confirmed by mapping observations.
(3) Comparison between the kinematics of the methanol
masers and of the thermal molecular lines reveals that they trace
a wide range of molecular cloud conditions. In some objects the
maser emission occurs in a narrow velocity range centered at the
systemic velocity, which may indicate that the innermost parts
of a molecular cloud or a circumstellar disc is the site of maser
emission. In other objects the velocities of maser features are
very similar to, or slightly overshoot, the velocity ranges of the
thermal molecular line wings, suggesting that the masers arise
in outflows. There are also objects where the maser emission re-
veals more complex kinematics.
(4) The column density of H2 derived from the CO and
HCO+ lines are between 1022 and 2 × 1023 cm−2. We use
our measurements of the intensity ratios of the CS lines to in-
fer that methanol masers arise from regions with a gas den-
sity of 105 − 107 cm−3, a kinetic temperature of 30 − 100K
and a methanol fractional abundance of 5 × 10−7 − 10−5. This
represents a significant refinement to the input parameters of
methanol maser models.
Acknowledgements. We like to thank the staff of the IRAM 30 m telescope for
help with the observations and the unknown referee for helpful comments. This
work has been supported by the Polish MNiI grant 1P03D02729.
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	Introduction
	The sample
	Astrometric positions
	Distances
	Observations and data reduction
	Results
	Systemic velocities
	Shape of profiles
	Asymmetry
	Wings
	Derivation of physical parameters
	Column densities
	Gas density and temperature
	Discussion
	Kinematics
	Implications for the evolutionary status 
	Constraints on maser models
	Conclusions
