A&A 460, 721-731 (2006)
DOI: 10.1051/0004-6361:20064815
O. Miettinen - J. Harju - L. K. Haikala - C. Pomrén
Observatory, P.O. Box 14, 00014 University of Helsinki, Finland
Received 5 January 2006 / Accepted 5 September 2006
Abstract
Aims. We determine the fractional SiO abundance in high-mass star-forming cores, and investigate its dependence on physical conditions, to provide constraints on the chemistry models of the formation of SiO in the gas phase or via grain mantle evaporation. The work addresses also CH3CCH chemistry, as the kinetic temperature is determined using this molecule.
Methods. We estimate the physical conditions of 15 high-mass star-forming cores and derive the fractional SiO and CH3CCH abundances using spectral line and dust continuum observations with the SEST.
Results. The kinetic temperatures as derived from CH3CCH range from 25 to 39 K, the average being 33 K. The average gas density in the cores is
cm-3. The SiO emission regions are extended and typically half of the integrated line emission comes from the velocity range traced out by CH3CCH emission. The upper limit of SiO abundance in this "quiescent'' gas component is
10-10. The average CH3CCH abundance is about
.
It shows a shallow, positive correlation with the temperature, whereas SiO shows the opposite tendency.
Conclusions. We suggest that the high CH3CCH abundance and its possible increase when the clouds become warmer is related to the intensified desorption of the chemical precursors of the molecule from grain surfaces. In contrast, the observed tendency of SiO does not support the idea that the evaporation of Si-containing species from the grain mantles would be important, and it contradicts models where neutral reactions with activation barriers dominate SiO production. A possible explanation for the decrease is that warmer cores represent more evolved stages of core evolution with fewer high-velocity shocks and thus less efficient SiO replenishment.
Key words: ISM: clouds - ISM: molecules - molecular data - radio continuum: ISM - radio lines: ISM - stars: formation
High-mass star-forming regions with large molecular column densities and bright molecular lines offer excellent opportunities to study interstellar chemisty. Results from chemistry models can be combined with observational data on the physical conditions and molecular abundances. This information can help us to link the properties of massive dense cores, their chemical characteristics, and phenomena like massive outflows, ultracompact (UC) H II regions and molecular masers into a coherent evolutionary sequence. This is useful since the process of high-mass star formation is not yet well understood (see, e.g., Garay & Lizano 1999; Fontani et al. 2002; Thompson & Macdonald 2003).
In this paper we derive the abundances of silicon monoxide, SiO, and methyl acetylene, CH3CCH, in a sample of massive molecular cloud cores, and discuss the relation of these abundances to the physical conditions and the probable evolutionary stages of the cores. The purpose is to contribute towards a better understanding of the interplay between the processes associated with massive star formation and the physical and chemical properties of the surrounding dense core.
SiO is believed to trace exclusively shocked gas as its spectral lines have wings and are often shifted relative to the emission from the ambient gas (see, e.g., Martín-Pintado et al. 1992; Schilke et al. 1997). Our current understanding is that SiO is evaporated from the dust grains when the shock velocity is greater than about 20 km s-1. This, however, depends on the grain-mantle composition in the preshock gas (e.g., Schilke et al. 1997). In dense and warm giant molecular cloud (GMC) cores, SiO seems to present also in the quiescent gas component. In the few interferometric maps available, extended SiO emission at ambient cloud velocity is either seen in the vicinity of high-velocity outflows (Hatchell et al. 2001) or totally separate from them (Lefloch et al. 1998; Codella et al. 1999; Shepherd et al. 2004; Fuente et al. 2005). The fractional abundance of "quiescent SiO'' is estimated to be about 10-10 in GMC cores. SiO absorption measurements towards dormant GMCs, so-called "spiral arm clouds'', give similar abundances (Greaves et al. 1996). The present observational evidence is insufficient to determine whether the presence of SiO is caused by thermal evaporation, enhanced neutral-neutral production pathways, photon-induced reactions, or by shock removal. To resolve this we determine fractional SiO abundances and kinetic temperatures in a representative sample of GMC cores, and correlate these also with the kinematic information contained in the spectral lines.
CH3CCH is an organic molecule observed frequently towards
dense cores. Its abundance is of the order of 10-9.
The rotational temperature,
,
derived from
a series of
JK-(J-1)K rotational lines
of CH3CCH is considered a good estimate of
the gas kinetic temperature,
,
in molecular clouds
(e.g., Askne et al. 1984; Bergin et al. 1994).
Since CH3CCH has a relatively low dipole
moment (
D, Bauer et al. 1979; Burrell et al. 1980),
its rotational levels are thermalized at densities
or higher
(e.g., Kuiper et al. 1984). This property together with
the assumption that CH3CCH emission is optically thin (
),
allows us to assume local thermodynamic equilibrium (LTE),
so that
may be derived. We also use
to
estimate the dust temperature,
,
which is needed to
derive the total gas column densities (
)) from dust
continuum observations.
The source list consists of 15 high-mass star-forming cores associated with OH, H2O, and CH3OH masers, UC H II regions and bright, thermal SiO rotational line emission. The sources were selected from the SiO survey of Harju et al. (1998) (hereafter HLBZ98). The positions selected for the SiO and CH3CCH line observations are listed in Table 2. These positions served as map centres in the continuum observations.
The SiO and CH3CCH column densities and the kinetic temperatures
in the present paper are derived from spectral line observations with the
SEST. The estimates for the molecular hydrogen column densities,
,
and the core masses are derived from dust continuum maps
obtained with the SIMBA bolometer at SEST. The observations and
the data reduction procedures are described in Sect. 2.
The direct observational results are presented in Sect. 3.
In Sect. 4 we describe the methods used to derive
the physical/chemical properties of the GMC cores. In Sect. 5
we discuss the results, and in Sect. 6 we summarize our major conclusions.
The spectral line observations were made
during four observing runs from 1995 to 2003 with
the Swedish-ESO-Submillimetre Telescope SEST at
the La Silla observatory, Chile. The SEST 3 and 2 mm (SESIS) dual SIS single sideband (SSB) receiver was used. The observations were made
in the dual beam switching mode (beam throw
in azimuth).
The SEST half-power beam width (HPBW) and the main beam efficiency,
,
at frequencies 86 GHz, 115 GHz and 147 GHz are
,
,
and 0.75, 0.70, 0.66, respectively. The
SEST high resolution 2000 channel acousto-optical spectrometer
(bandwidth 86 MHz, channel width 43 kHz) was split into two halves to
measure two receivers simultaneously. At the observed wavelengths, 2 mm
and 1 mm, the 43 kHz channel width corresponds to approx. 0.12 km s-1 and 0.08 km s-1, respectively. The observed molecular transitions, their rest
frequencies and the upper level energies are listed in Table 1.
Calibration was achieved by the chopper wheel method. Pointing was
checked regularly towards known circumstellar SiO masers.
Pointing accuracy is estimated to be better than
.
The 28SiO (v=0, J=2-1) and (v=0, J=3-2) observations were carried out in October 1995 and April 1996 and are described in detail in HLBZ98. 29SiO(2-1) and (3-2) observations were carried out in October 1998. Typical values for the effective SSB system temperatures and the rms noise of the spectra were 140 K and 170 K, and 0.016 K and 0.021 K at frequencies 86 GHz and 130 GHz, respectively. Linear baselines were subtracted from the spectra.
The CH3CCH ( J=5K-4K) and ( J=6K-5K) observations were carried out in June 2003. At this time the electronics of the SESIS 2 mm receiver had been decommissioned and only the 3 mm receiver was available. These were the last spectral line observations made with SEST. For each source typically six spectra, each with 2 min integration time, were obtained. The effective SSB system temperatures were 130 K for the CH3CCH(5K-4K) transition and 250 K for the CH3CCH(6K-5K). The rms of the spectra were 0.017 K and 0.024 K for the two transitions, respectively. Because of the better performance of the receiver at the lower frequency, the J=5K-4K transition was chosen for the survey. A third order baseline was subtracted from the spectra. Four K-components (K=0,1,2,3) were detected towards all sources, and Gaussian profiles were fitted to the lines. The CLASS program of the GILDAS software package developed at IRAM was used for the reductions. In the fitting procedure the line separations between the components were fixed and the line widths for all K-components were assumed to be identical.
The 1.2 mm continuum observations were carried out in June 2003
with the 37 channel SEST imaging bolometer array, SIMBA.
The SIMBA central frequency is 250 GHz and the bandwidth about 50 GHz.
The HPBW of a single bolometer element is
and the separation between elements on the sky is
.
The observations were conducted in stable weather.
Frequent skydips were used to
determine the atmospheric opacity and the values obtained varied
between 0.27 and 0.3. Pointing was controlled by observing sources
with known accurate coordinates and is estimated to be better than 5
.
Uranus was used for flux density calibration.
The observations were done in the fastscanning mode with a scanning
speed of
.
The typical map consisted
of 65 scans of
in length in azimuth and spaced by
in elevation. Each source was observed at least twice. The two maps were made
at different LST values to reduce the possibility of strong artefacts
in the reduced maps. The data were reduced using the
MOPSI
software package according to
guidelines in the SIMBA Observers Handbook (2002) and Chini et al. (2003).
All the observed maps contained strong sources and therefore, as suggested in the SEST handbook, no despiking was applied to the data. The data reduction deviated from the suggestions in the SEST handbook in the following details. The SEST dish shape (and the beam profile) degrades as a function of the elevation as the dish deforms because of its own weight. Gain-elevation correction is used to compensate the decrease of the observed source peak intensity as a function of the telescope elevation. However, no gain-elevation correction was applied to the present data because aperture photometry was used to define the source flux. The source flux lost because of the decrease of the telescope gain in the source position is recovered as the source is mapped. The telescope gain is unity at an elevation of 74.6 degrees and decreases by 10% at 45 degrees. Most of the present observations were obtained at elevations of 50 degrees or higher where the said correction is small. A third order baseline was applied to the data before correlated noise removal after which all the bolometer elements with insufficient base range were masked. A first order baseline was applied to individual scan lines after this. The refinement of the SIMBA data reduction is discussed in detail in Haikala (2006, in preparation).
The observed sources are listed in Table 2.
The columns of this table are: (1) source name; (2) and
(3) equatorial coordinates (J2000.0); (4) galactic
coordinates; (5) distance; (6) galactocentric distance;
(7) LSR velocity determined from the CH3CCH line;
(8) notes on association with molecular masers or an
UC H II region. When no distance
reference is given in Col. 5, the value is the kinematic distance
calculated from the CH3CCH velocity using the rotation curve determined by
Brand & Blitz (1993) and R0=8.5 kpc (galactocentric distance of the Sun)
and
km s-1(circular velocity at a distance R0).
The obtained SIMBA maps are presented in Fig. 1.
The maps are plotted on the same angular and intensity scales.
In Table 3 we list the coordinates
(
,
)
and
the intensities (
)
at the dust emission maxima identified on the maps.
The SEST HPBW at the frequency of the continuum observations is 24
which is less than half of the HPBW at the frequency of the
observed SiO J=2-1 and the CH3CCH lines. In order to estimate
the average dust and H2 column densities within
the beam used for the 3 mm line observations, the SIMBA maps were
smoothed to the resolution of 57
.
The resulting smoothed surface
brightness (
)
towards the position of the line measurements
is listed in Col. 2 of Table 4.
The 1.2 mm flux density (
)
integrated over the source and
the angular size (FWHM) of the core (
)
determined
from the original map is listed in Cols. 3 and 4 of this table.
The angular sizes were estimated from
two-dimensional Gaussian fits to the surface brightness
distribution (geometric mean of the observed major and
minor axes of Gaussians).
The source size was corrected for the broadening by 24
the telescope beam in accordance with the assumption of Gaussian
distributions. The flux density was calculated in 72
diameter
circular aperture which correponds to three times the original beam size.
Source | ![]() |
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IRAS 12326-6245 | 134 | 12.8 | 13 |
G326.64+0.61 | 98 | 10.9 | 27 |
OH328.81+0.63 | 136 | 14.6 | 22 |
IRAS 15566-5304 | 74 | 7.8 | 34 |
G330.95-0.19 | 154 | 22.8 | 13 |
G345.01+1.8N | 154 | 16.4 | 47 |
IRAS 16562-3959 | 185 | 19.1 | 36 |
G345.00-0.23 | 110 | 11.4 | 19 |
NGC 6334 FIR-V | 189 | 19.4 | 39 |
NGC 6334F | 360 | 34.7 | 16 |
G351.77-0.54 | 289 | 27.5 | 17 |
G353.41-0.36 | 197 | 22.6 | 33 |
W28 A2(2) | 163 | 15.5 | 15 |
W31 (2) | 198 | 19.1 | 89 |
W33 CONT | 395 | 41.1 | 32 |
The obtained SiO spectra are presented in Fig. 2.
The lines are nearly Gaussian in the central part and have broad,
asymmetric wings.
The line parameters are listed in Table 6.
In this table we give the minimum, maximum, and peak velocities
(
,
and
,
respectively),
and the integrated intensities,
,
over the given velocity range.
The spectra of the two SiO isotopologues were used to estimate
the optical thicknesses of the J=2-1 and J=3-2 lines
of 28SiO,
and
.
For this purpose all spectra were resampled to the
same LSR velocity grid with a channel width of 1 km s-1.
The smoothing of the spectra to this resolution was necessary to achieve
a sufficient signal to noise ratio. Using these resampled
spectra the optical thicknesses
and
could be derived in 1-13 velocity channels near
the line peak for twelve sources. The terrestrial SiO isotopic abundance
ratio is
.
In the interstellar
medium, values in the range 10-20 have been determined (Penzias 1981).
Here we use the value X=20 when determining the optical thicknesses.
The 28SiO optical thicknesses
and
are typically between 1 and 3 in the line
centres.
The excitation temperatures,
,
were estimated
from the optical thickness ratio by using Eq. (2) of HLBZ98. These estimates are based on the assumption that
.
They do not depend on the
beam filling as the optical thicknesses were determined using pairs
of lines having similar frequencies. The derived values of
are
close to 5 K for all sources and velocity channels.
The weighted average of
and its standard deviation are
K (see Table 5).
From these results, the SiO column densities, N(SiO), were
estimated from the integrated intensities of the 29SiO spectra
by assuming optically thin emission and uniform excitation of
the rotational lines with
K for all sources.
The CH3CCH spectra are shown in Fig. 3,
and the Gaussian parameters of the detected lines
are given in Table 7.
The line peak velocities and
widths (FWHM) are listed in Cols. 2 and 3 of this table. The integrated
intensities of the components K=0,1,2 and 3 are given in Cols. 4-7.
The integrated intensities were used to derive the rotational
temperatures,
,
and the CH3CCH column densities,
,
with the population diagram method described and
discussed in, e.g., Askne et al. (1984), Bergin et al. (1994), and
Goldsmith & Langer (1999).
In the case of uniform excitation and optically thin dipole transition,
the total column density of the emitting molecule is related
to the integrated line intensity by the formula
The rotational temperatures,
,
and the CH3CCH column densities,
,
were derived by means of the population diagram method.
This method is based on a comparison of intensities of spectral lines
that lie close in frequency although arising from rotational levels with
different energies. In the case of CH3CCH, the close lying transitions
represent different K-components.
Assuming uniform excitation and optically thin emission,
the "rotational diagram equation'', i.e. the equation relating
the integrated intensities of different K-components to
the rotational temperature,
,
and
the total column density,
,
can be written as
Rotational diagrams for two sources are shown in
Fig. 4.
Straight lines were fitted to the data using a least-squares fitting technique.
The data are consistent with a single
.
The goodness of
the fit substantiates the assumptions that the gas is in LTE with
and that the lines are optically
thin.
Source |
![]() |
![]() |
![]() |
![]() |
G326.64+0.61 | 3.0 ![]() |
2.5 ![]() |
5.2 ![]() |
2 |
OH 328.81+0.63 | 2.1 ![]() |
2.1 ![]() |
4.8 ![]() |
10 |
IRAS 15566-5304 | 3.8 ![]() |
2.6 ![]() |
4.7 ![]() |
5 |
G345.01+1.8N | 1.5 ![]() |
1.3 ![]() |
5.3 ![]() |
1 |
G345.00-0.23 | 2.2 ![]() |
1.5 ![]() |
4.6 ![]() |
6 |
NGC 6334 FIR-V | 1.7 ![]() |
1.1 ![]() |
4.1 ![]() |
1 |
NGC 6334F | 2.3 ![]() |
1.2 ![]() |
3.4 ![]() |
3 |
G351.78-0.54 | 1.9 ![]() |
1.7 ![]() |
4.2 ![]() |
13 |
G353.41-0.36 | 2.2 ![]() |
2.0 ![]() |
6.1 ![]() |
5 |
W28 A2(2) | 2.4 ![]() |
2.0 ![]() |
5.2 ![]() |
3 |
W31 (2) | 2.9 ![]() |
2.9 ![]() |
4.0 ![]() |
6 |
W33 CONT | 2.7 ![]() |
3.0 ![]() |
7.2 ![]() |
1 |
The H2 column density, N(H2), of each source
was derived from dust continuum emission using the
following equation:
We assume that
is equal to the gas kinetic
temperature derived from CH3CCH observations.
and
are believed to be well coupled deep
in the cloud, above a density of about 104 cm-3,
due to collisions between gas and dust grains
(see, e.g., Takahashi et al. 1983; Lee et al. 2004).
Values of 0.1 m2 kg-1 at
mm (Ossenkopf & Henning 1994)
and
are adopted for
and
,
respectively.
The rotational temperatures, the H2, SiO and CH3CCH column densities and
the fractional SiO and CH3CCH abundances
(
)
are listed
in Table 8.
The average value and the standard deviation of
is
K. The N(H2) values are found to be
1022-1023 cm-2, and N(SiO) and N(CH3CCH) lie
in the range
cm-2 and
cm-2, respectively.
The values of
and
are
found to be
and
,
respectively.
Line centre | Whole line | |||||||
Source | Line |
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[km s-1] | [km s-1] | [km s-1] | [K km s-1] | [km s-1] | [km s-1] | [K km s-1] | ||
IRAS 12326-6245 | J=2-1 | -41.8 | -37.2 | -38.4 | 0.10(0.01) | -49.3 | -35.0 | 0.20(0.02) |
J=3-2 | -40.0 | 0.23(0.02) | -48.4 | -32.0 | 0.40(0.03) | |||
G326.64+0.61 | J=2-1 | -42.2 | -37.7 | -40.7 | 0.26(0.01) | -46.1 | -33.6 | 0.47(0.02) |
J=3-2 | -41.0 | 0.17(0.02) | -45.5 | -33.6 | 0.31(0.03) | |||
OH328.81+0.63 | J=2-1 | -44.6 | -38.3 | -41.9 | 0.51(0.01) | -55.3 | -24.9 | 0.99(0.03) |
J=3-2 | -42.5 | 0.60(0.02) | -49.6 | -31.4 | 1.07(0.03) | |||
IRAS 15566-5304 | J=2-1 | -46.2 | -41.0 | -45.8 | 0.41(0.01) | -53.1 | -33.4 | 0.85(0.03) |
J=3-2 | -44.6 | 0.26(0.02) | -50.1 | -36.7 | 0.50(0.03) | |||
G330.95-0.19 | J=2-1 | -92.7 | -88.7 | -91.6 | 0.16(0.01) | -99.6 | -80.5 | 0.37(0.03) |
J=3-2 | -88.4 | 0.13(0.02) | -97.5 | -85.0 | 0.20(0.03) | |||
G345.01+1.8N | J=2-1 | -16.8 | -10.6 | -16.4 | 0.31(0.02) | -23.5 | -6.5 | 0.50(0.03) |
J=3-2 | -17.3 | 0.24(0.02) | -21.1 | -7.4 | 0.39(0.03) | |||
IRAS 16562-3959 | J=2-1 | -14.9 | -8.7 | -12.2 | 0.21(0.02) | -19.1 | -0.7 | 0.35(0.03) |
J=3-2 | -11.3 | 0.27(0.02) | -21.8 | -2.4 | 0.62(0.03) | |||
G345.00-0.23 | J=2-1 | -29.7 | -22.4 | -28.2 | 0.43(0.02) | -40.4 | -13.9 | 0.88(0.03) |
J=3-2 | -26.9 | 0.48(0.02) | -40.7 | -11.2 | 1.09(0.04) | |||
NGC 6334 FIR-V | J=2-1 | -9.3 | -2.9 | -6.7 | 0.23(0.02) | -20.7 | -3.7 | 0.51(0.03) |
J=3-2 | -7.0 | 0.29(0.02) | -17.7 | 1.4 | 0.56(0.03) | |||
NGC 6334F | J=2-1 | -8.5 | -2.7 | -8.8 | 0.22(0.01) | -11.1 | -2.5 | 0.34(0.01) |
J=3-2 | -10.1 | 0.20(0.01) | -13.8 | -0.1 | 0.37(0.02) | |||
G351.77-0.54 | J=2-1 | -5.1 | 0.8 | -3.9 | 0.72(0.01) | -19.0 | 11.0 | 1.85(0.03) |
J=3-2 | -3.7 | 0.91(0.02) | -17.8 | 11.0 | 2.29(0.03) | |||
G353.41-0.36 | J=2-1 | -19.7 | -13.6 | -16.7 | 0.39(0.01) | -24.7 | -6.2 | 0.66(0.02) |
J=3-2 | -17.1 | 0.30(0.01) | -21.7 | -9.2 | 0.42(0.02) | |||
W28 A2(2) | J=2-1 | 6.5 | 12.8 | 10.2 | 0.23(0.02) | -2.5 | 26.6 | 0.69(0.04) |
J=3-2 | 11.9 | 0.35(0.02) | 0.1 | 37.9 | 1.34(0.05) | |||
W31 (2) | J=2-1 | -5.2 | 1.6 | -4.5 | 0.39(0.02) | -12.8 | 7.2 | 0.68(0.03) |
J=3-2 | -2.1 | 0.41(0.02) | -12.5 | 5.7 | 0.67(0.03) | |||
W33 CONT | J=2-1 | 31.9 | 39.2 | 35.0 | 0.24(0.02) | 28.9 | 48.0 | 0.33(0.03) |
J=3-2 | 35.4 | 0.28(0.02) | 27.1 | 47.7 | 0.48(0.03) |
The linear sizes have been computed from the angular diameters listed in Table 4 using the distances listed in Table 2.
The core masses have been estimated using the 1.2 mm continuum maps.
We have also calculated the virial masses using the velocity
dispersions and kinetic temperatures from the CH3CCH data.
The mass of a core,
,
has been calculated assuming that
the dust emission is optically thin and
that the dust-to-gas ratio,
,
and the dust absorption coefficient
per unit mass,
,
are constant. These assumptions imply the following equation:
The virial masses,
,
have been estimated by approximating
the mass distribution by a homogenous, isothermal sphere without magnetic
support and external pressure, using the formula
The average H2 number densities,
,
were calculated using the
masses,
,
and radii, R, estimated from the dust
continuum maps. The obtained radii, masses and average densities are listed
in Table 9. The average radius of the sources
is
pc, the average value of the masses estimated
from continuum emission is
,
and the average density is
cm-3.
![]() |
Figure 4:
Rotational diagrams of the CH3CCH spectra for two of the observed
sources. X-axis plots the upper transition state energies divided by
the Boltzmann constant,
![]() ![]() |
The 1.2 mm dust emission is likely to be dominated by the cool, K, envelopes, although it has a contribution from hot cores (T > 100 K)
around newly born massive stars. The dust temperatures derived by
Faúndez et al. (2004) towards several of our sources using SIMBA and IRAS 100 and 60
m data are comparable to the CH3CCH rotational
temperatures,
,
derived here. Taking the uncertainties due
to different beam sizes of the SIMBA and IRAS into account, the
agreement is reasonable and substantiates the assumption
used in the total
estimates.
While it can be assumed that the CH3CCH lines and 1.2 mm dust continuum emission trace to a large part the same material, this is not obvious for dust and SiO. All SiO lines observed in this survey show high-velocity wings, are single-peaked, and have their maxima near the systemic velocity of the cloud. These characteristics can be qualitatively explained with bow-shock models where the jet generating the shock is seen at a small angle, i.e. either head on or tail on. A large inclination should result in a double peaked profile from a bow shock, and in a narrow line in the case where the emission originates in a turbulent wake behind a bow shock (HLBZ98, Fig. 12). In view of the distance to the sources and the angular resolution of the present observations, it is clear that the telescope beam encompasses entire star-forming regions with various gas components and young stars at different evolutionary stages. The line profiles may therefore have a contribution from several outflows with different inclinations.
In the light of previous SiO mapping observations it seems possible, however,
that part of the low-velocity emission originates in the quiescent envelope.
This kind of component has been discovered in
high-resolution SiO mappings of star-forming regions
(Lefloch et al. 1998; Codella et al. 1999; Shepherd et al. 2004;
Fuente et al. 2005). In these studies SiO emission characterized by
narrow and broad lines have been found to come from separate
regions. The SiO abundances derived for the low-velocity component
are 10-10. This value is smaller than the abundances
derived towards high-velocity shocks (
10-8, e.g.,
Bachiller et al. 1991; Gibb et al. 2004), and larger than the upper
limits derived in cold, dark clouds and PDRs (
10-12, e.g.,
Ziurys et al. 1989; Martín-Pintado et al. 1992; Schilke et al. 2001).
In the present single-pointing observations typically half of the integrated SiO intensity comes from the velocity range with detectable CH3CCH, which is supposed to trace the quiescent gas component (see Table 6). However, the same radial velocity does not necessarily mean spatial coincidence. Although SiO abundances derived in outflows are based on comparison with CO in the same velocity range, a similar spatial velocity correlation cannot be assumed near the cloud's systemic velocity. Depending on the structure and orientation of the shock, it is possible to find shock-produced SiO at low radial velocities. Therefore the column density of the low-velocity SiO gives only an upper limit to the SiO column density in the quiesent envelope, and the same is probably true for the fractional SiO abundance using the H2 column density determined by dust.
A comparison of the SiO main beam brightness temperatures,
,
with the derived optical thicknesses,
,
and excitation
temperatures,
,
suggests that the beam filling factors
are close to unity, and the SiO emission regions are extended. We find
that the upper limits for the fractional SiO abundance in the
low-velocity gas are of the order of 10-10, i.e. similar to the
SiO abundances derived previously for the quiescent gas component in some
star-forming regions. It should be noted, however, that the same
result would be achieved if SiO is present only in shocked layers
filling 1% of the gas volume. This is not necessarily in
contradiction with nearly uniform beam filling.
The model for greatly enhanced SiO production in powerful shocks by
silicate grain destruction and subsequent high-temperature gas-phase
chemistry is well established (Schilke et al. 1997; Caselli et al. 1997;
Pineau des Forêts et al. 1997). The possible occurrence of SiO in the quiescent
gas component is not as well understood. In the postshock gas SiO in
a reaction with OH should convert into SiO2 (SiO + OH
SiO2 + H), which is eventually removed from the
gas phase by accretion onto the dust grains.
According to Codella et al. (1999), the times needed to remove SiO from
the gas phase and to slow down speeding SiO "blobs'' are similar,
104 yr, and this would support the outflow remnant scenario. In
another, perhaps more feasible model high-velocity SiO emission comes
from bow-shocks generated by high-velocity jets. In this case
low-velocity SiO can originate in turbulent wakes behind
bow-shocks (see Raga & Cabrit 1993, and the model profiles in Fig. 12 of HLBZ98).
Source |
![]() |
![]() |
![]() |
![]() |
![]() |
![]() |
[K] | [
![]() |
[
![]() |
[10-9] | [
![]() |
[10-10] | |
IRAS 12326-6245a | 34.0 | 6.4 | 2.0 | 3.1 | 0.8 | 1.3 |
G326.64+0.61 | 29.0 | 5.7 | 3.4 | 6.0 | 2.0 | 3.5 |
OH328.81+0.63 | 30.5 | 7.4 | 4.9 | 6.6 | 3.8 | 5.1 |
IRAS 15566-5304 | 27.2 | 4.7 | 2.2 | 4.7 | 3.1 | 6.6 |
G330.95-0.19 | 31.9 | 8.0 | 3.5 | 4.4 | 1.2 | 1.5 |
G345.01+1.8N | 34.5 | 7.3 | 3.5 | 4.8 | 2.3 | 3.2 |
IRAS 16562-3959 | 35.9 | 8.3 | 15.0 | 18.0 | 1.6 | 1.9 |
G345.00-0.23 | 30.6 | 6.0 | 4.0 | 6.7 | 3.3 | 5.5 |
NGC 6334 FIR-V | 33.6 | 9.2 | 5.8 | 6.3 | 1.8 | 2.0 |
NGC 6334F | 35.8 | 16.3 | 6.3 | 3.9 | 1.6 | 1.0 |
G351.77-0.54 | 34.4 | 13.7 | 8.9 | 6.5 | 5.5 | 4.0 |
G353.41-0.36 | 25.2 | 13.7 | 4.2 | 3.1 | 3.0 | 2.2 |
W28 A2(2) | 34.6 | 7.7 | 11.2 | 14.6 | 1.7 | 2.2 |
W31 (2) | 35.0 | 9.2 | 6.2 | 6.8 | 2.9 | 3.2 |
W33 CONT | 38.6 | 16.4 | 14.3 | 8.7 | 1.8 | 1.1 |
a The rotational temperature and CH3CCH column density derived from the J=6K-5K line are 37.2 K and ![]() |
It has been suggested that SiO can be produced also in warm, quiescent
gas via neutral or ion-neutral reactions, possibly preceded by
evaporation of Si-bearing molecules from the icy mantles of dust
grains. The suggestion of Ziurys et al. (1989) and Langer & Glassgold (1990)
that neutral reactions with activation energies are important for the
SiO production in warm gas would imply a strong correlation between
the SiO abundance and the average kinetic temperature. As shown in
Fig. 5, bottom panel, even in the rather
narrow temperature range covered, 25-39 K, a substantial change in
the fractional SiO abundance would be expected if it were proportional
to
as suggested by Langer & Glassgold (1990). In the
model of MacKay (1995, 1996) icy silicon is
mainly in the form of SiH4, which can evaporate when the dust is
warmed up. Also in this model one could expect the warmer cores to have
larger SiO abundances.
The abundances derived for the low-velocity SiO do not show any positive correlation to the average kinetic temperature. Instead, they even seem to decrease slightly when the temperature rises (Fig. 5, bottom panel).
The CH3CCH production can be enhanced by intensified
desorption. CH3CCH is one of the lowest energy products in the
reaction between the methylidine radical (CH) and ethylene
(C2H4, Canosa et al. 1997) which is formed on dust grains. The
diagram presented in Fig. 5, top panel,
suggests a positive correlation between the fractional CH3CCH
abundance and
,
which would conform with the desorption
scenario. On the other hand, this tendency does not agree with the
prediction of the pure gas-phase model of Lee et al. (1996; see also
Alakoz et al. 2002), which predict a strong anticorrelation between the
CH3CCH abundance and the kinetic temperature. The CH3CCH
abundances derived here are similar to those found in the Orion
Ridge, M17 and Cepheus A (Ungerechts et al. 1997; Bergin et al. 1997),
and substantially larger than those determined towards cold, starless
cores (Kontinen et al. 2000; Markwick et al. 2005).
These findings lend support to the possibility that grain-surface reactions are important for the production of CH3CCH and its gas-phase abundance depends on temperature via the thermal desorption. On the other hand, no evidence for such behaviour is found in the case of SiO. This complies with the shock origin of SiO. The diminishing tendency towards warmer cores suggested by Fig. 5 may reflect an evolutionary effect.
According to the picture presented by Fuente et al. (2005) based on a
study of the environments of two young stars in the region of the
nebula NGC 7129, protostellar envelopes are dispersed and warmed up
during the early stellar evolution, at the same time as shocks
associated with outflows become less energetic. For SiO this means
decreasing abundance with time. The abundance variation from core to
core is not very large in our sample. This is probably related to the
fact that all the objects studied are associated with powerful masers
and/or UC H II regions, and therefore the dynamical
and chemical age variations are likely to be rather limited
(Garay & Lizano 1999; Bergin et al. 1997).
Source | R |
![]() |
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[pc] | [![]() |
[![]() |
[cm-3] | |
IRAS 12326-6245 | 0.1 | 2200 | 410 |
![]() |
G326.64+0.61 | 0.2 | 850 | 430 |
![]() |
OH328.81+0.63 | 0.2 | 1200 | 320 |
![]() |
IRAS 15566-5304 | 0.2 | 810 | 740 |
![]() |
G330.95-0.19 | 0.2 | 8100 | 1200 |
![]() |
G345.01+1.8N | 0.1 | 410 | 550 |
![]() |
IRAS 16562-3959 | 0.1 | 400 | 400 |
![]() |
G345.00-0.23 | 0.1 | 960 | 870 |
![]() |
NGC 6334 FIR-V | 0.2 | 500 | 420 |
![]() |
NGC 6334F | 0.1 | 830 | 360 |
![]() |
G351.77-0.54 | 0.03 | 120 | 150 |
![]() |
G353.41-0.36 | 0.3 | 3500 | 960 |
![]() |
W28 A2(2) | 0.1 | 970 | - |
![]() |
W31 (2) | 0.7 | 1600 | 3500 |
![]() |
W33 CONT | 0.3 | 4300 | 900 |
![]() |
![]() |
Figure 5:
Fractional CH3CCH and SiO abundances as functions
of the rotational temperature,
![]() ![]() |
Most of the cores studied here appear to be gravitationally bound
which is indicated by the fact that
.
The
ratios range from 0.7 to 6.8, the
average being 2.5. For five cores
is substantially
larger than
,
whereas the rest appear to be close to the
virial equilibrium (
,
within a factor of two).
The virial masses,
,
were
calculated assuming homogenous density distributions. A more
realistic distribution described by a power law of the
type
,
where p is
typically
1.5-2.5, would make the
virial mass estimates smaller (see, e.g., Fontani et al. 2005).
On the other hand, the effect of the
density gradient can be compensated by the increase of the kinetic
temperature towards the core centre. All the cores studied here
contain newly born massive stars which are known to give rise to
expanding UC H II regions and hot cores.
Magnetic fields are likely to provide
additional support against gravitational collapse (e.g., Garay & Lizano 1999;
Fontani et al. 2002).
For example, Sarma et al. (2000) detected magnetic fields toward sources A, E, and D in the NGC 6334 complex, which are of the same order as
the critical fields required to support the cores against
gravitational collapse.
Ten of our fourteen targets have been included in recently published, extensive millimetre continuum surveys with SIMBA. Faúndez et al. (2004) have mapped about 150 massive cores around IRAS point sources with FIR colors typical of UC H II regions, and discuss the general characteristics of these objects. They present the maps of IRAS 12326-6245, G326.64+0.61 (IRAS 15048-5356), OH328.81+0.63 (IRAS 15520-5234), G330.95-0.19 (IRAS 16060-5146), IRAS 16562-3959, G345.00-0.23 (IRAS 17016-4124), NGC 6334F (IRAS 17175-3544), and G351.77-0.54 (IRAS 17233-3606). The survey of Hill et al. (2005) covering about 130 regions including IRAS 12326-6245 (G301.14-0.2), G330.95-0.18, W28 A2 (G5.89-0.39), and W31 (2) (G10.62-0.38).
The structure of the sources in Fig. 1
agrees well with those presented in Faúndez et al. (2004) and Hill et al. (2005).
The peak fluxes in Hill et al. (2005) for the sources in common with this work
also agree well. The integrated flux densities for the sources are however
higher in Hill et al. (2005) and especially in Faúndez et al. (2004) than
in the present paper. Hill et al. calculated the integrated flux
density inside a polygon containing all the emission around the source
whereas in the present paper the integration was done inside a circle
corresponding to three times the telescope beam width.
Contrary to the present paper Hill et al. (2005) also applied the gain-elevation correction to the SIMBA data. The different integration aperture
and the gain-elevation correction can well explain the Hill et al. (2005)
higher flux densities. Faúndez et al. (2004) give no details on
how the integrated fluxes were calculated so it cannot be decided
if the higher values are due to e.g. the integration aperture.
The kinetic temperatures from CH3CCH are lower than
the dust temperatures derived by Faúndez et al. (2004) from the spectral
energy distributions (SEDs) using the flux densities at 1.2 mm and in
the four IRAS bands. However, as our 1.2 mm flux densities are, in general,
lower than those derived by Faúndez et al.,
we end up with comparable masses.
The masses derived by Hill et al. (2005) are on average larger
than those derived by us, mainly because they have obtained
slightly higher flux densities and assumed that
K.
Finally, the mass estimates are proportional to the distance squared, and
there are some differences in the derived kinematic distances depending on the
Galactic rotation curve and
used.
Next we discuss six of our sources, IRAS 12326-6245, NGC 6334, G351.77-0.54, W28 A2(2), W31 and W33 in more detail. In general, the sources are typical high-mass star-forming cores associated with several molecular masers, UC H II regions and infrared sources (see Table 2 and Fig. 1).
IRAS 12326-6245 is a luminous FIR source (
)
which is located at a kinematically estimated distance
of 4.4 kpc (Zinchenko et al. 1995; Osterloh et al. 1997).
This object is associated with two UC H II regions,
both of which are associated with mid-infrared sources (Henning et al. 2000).
Also several molecular masers including H2O, OH and CH3OH have been
identified at the position of IRAS 12326-6245
(see Henning et al. 2000, for references).
The molecular line maps (see Henning et al. 2000, and references therein)
show that one of the most energetic and massive bipolar molecular outflow
(mass outflow rate
yr-1)
in the southern sky originates close to IRAS 12326-6245.
The core 1.3 mm flux density and the gas mass obtained by
Henning et al. (12.0 Jy and 2400 ,
respectively)
are very close to the values in the present study
(12.8 Jy and 2200
,
respectively).
The star-forming molecular ridge associated with the giant H II region NGC 6334 is the best studied of our targets. For recent, comprehensive molecular line and continuum studies see Kraemer et al. (1999), Kraemer & Jackson (1999), Jackson & Kraemer (1999), Sandell (2000), McCutcheon et al. (2000) and Brooks & Whiteoak (2001). These authors summarize the accumulated knowledge of the region, and include keys to the nomenclature. Sandell (2000) mapped the northern part of the molecular ridge with the UKT14 bolometer at the JCMT using four millimetre and submm wavelengths in the range 1.1-0.35 mm.
The six brightest 1.2 mm peaks listed in Table 3 can be associated with star-forming cores detected as far-infrared sources or UC H II regions as follows: 1) FIR-I/UC H II F; 2) I(N); 3) FIR-IV/UC H II A; 4) FIR-V; 5) FIR-III/UC H II C; 6) FIR-II/UC H II D (McBreen et al. 1979; Rodríguez et al. 1982; Gezari 1982). The 1.2 mm peak No. 7 is located on the southeastern side of the radio shell between FIR-IV and V, in the region of the PDR G351.2+0.70 (Jackson & Kraemer 1999). It is associated with a CO clump in the survey of Kraemer & Jackson (1999) (their Fig. 5). The dust emission peak No. 8 on the northwestern side of the dense molecular ridge is associated with the object called a "dust-cloudlet'' by Sandell (2000).
Assuming a uniform dust temperature of 30 K, we obtain from the SIMBA
map a mass of 11 400
for the entire ridge. The share of its
northern end with the cores I, I(N) and II is 5400
,
and for
the regions III, IV, and V we get 900
,
2200
,
and 1300
,
respectively.
The main features of the 1.1 mm map of Sandell (2000) can be recognized on the present SIMBA map, despite a lower S/N available here. A careful inspection of the image reveals filaments reaching out from the dense ridge, similar to the structures associated with OMC-1 discovered by Johnstone & Bally (1999) on a JCMT/SCUBA map. The suggestion of Johnstone & Bally (1999) that these features could represent the cavity walls of past outflows is appropriate also in this case. Filamentary structures are rather common features in GMCs forming high mass clusters, so there is no reason to consider them as artefacts of the data reduction.
The "linear filament'' discussed by Sandell (2000) continues to the
north and southwest and in fact the ridge is arc-shaped. In
Fig. 6 an arc of a circle of radius
(13 pc) and centred at RA
,
Dec
is superposed on the SIMBA map. The
alignment evident in the figure suggests that the ridge has been
formed by an expanding H II region with its centre
near the given position. This H II region has
probably been ionized by one or more of
the mature OB stars in the region. The closest O-type star on the
western side of the ridge is HD 319 699 (O7, distance 1.6 kpc).
The UV radiation from this star is likely to dominate the excitation of the
extended emission nebula in this direction. It seems, however,
that the star has not been at the centre of the expansion,
because it lies about
southeast of the centre of the circle,
and its proper motion vector (although very uncertain) points northwest
(the Tycho-2 Catalogue; Høg et al. 2000). The O-stars lying southeast and
south of the ridge, HD 319 702 (O9, d=1.4 kpc), HD 319 703
(O6, d=2.0 kpc), and HD 156 738 (O7, d=1.3 kpc) are surrounded by
separate, roundish ionization regions which probably have not contributed to
shaping the ridge (see, e.g., Fig. 5 of Kraemer & Jackson 1999 or
SkyView, http://skyview.gsfc.nasa.gov/).
G351.77-0.54 is the brightest SiO line source of the sample
(see Fig. 2).
It is associated with the luminous steep-spectrum far-infared source
IRAS 17233-3606, the strongest known ground-state OH maser source at
1665 MHz, other OH masers, and H2O and CH3OH masers (e.g.,
Caswell & Vaile 1995; MacLeod et al. 1998; Val'tts et al. 2000).
The kinematic distance derived using the LSR velocity from
CH3CCH is 700 pc, which would make this source the nearest
in our sample. The distance is smaller than the usually
adopted value (1.5-2.2 kpc), but it is close to the distance
derived by MacLeod et al. (1998) (1.0 kpc) from the LSR velocity of the
H2CO
110-111 absorption. The neighbourhood of the source
direction to that of the Galactic centre causes a large uncertainty in the
kinematic distance estimate. Nevertheless, the proximity is supported by the
intensity of the SiO emission, suggesting a large beam-filling factor.
G5.89-0.39 (also known as W28 A2) is one of the best-studied UC H II regions, but it probably is not the best understood. W28 A2 is associated with a highly energetic outflow, one of the most powerful ones in the Galaxy. In W28 A2 there may be multiple outflows and driving sources.
A bipolar outflow in the east-west direction with a mass of 70
was identified by Harvey & Forveille (1988)
in the J=1-0 lines of CO, 13CO, and C18O toward this source.
C34S observations shows a north-south oriented outflow
(Cesaroni et al. 1991). There is growing evidence that the O5 star in G5.89
produced the N-S molecular outflow and hence is forming through an
accretion process (see Arce et al. 2006 and references therein).
The SiO (v=0, J=1-0) observations made with the VLA by Acord et al. (1997)
show a northeast-southwest bipolarity, which is confirmed by
1.3 mm continuum and several molecular line observations by
Sollins et al. (2004). The 1.3 mm continuum source peaks
in the centre of the H II region (Sollins et al. 2004).
Moreover, NIR observations of the H2 v=1-0 S(1) rovibrational line and
the Br
emission show evidence of three outflows with
distinct orientations and driving sources (Puga et al. 2006).
Puga et al. suggest that the northwest-southeast oriented outflow is possibly
connected to the H2 knots seen in that direction.
One of them ("knot B''), NW of W28 A2, can be identified in our SIMBA map.
The two other features seen in our map, NW of knot B, at
,
,
and
,
,
could be related to NW-SE oriented outflow.
Thus it is unlikely that there exist a single central high-mass star in W28 A2 as was initially thought to be the case due to the symmetric morphology of the H II region.
The kinematic distance derived from H I observations towards
the supernova remnant W28 by Velázquez et al. (2002) is 1.9 kpc.
Adopting this value instead of the 2.7 kpc derived by us, the mass of the
cloud would decrease from 970
to 480
.
Thompson & Macdonald (1999) found a rotational temperature of 70 K
from CH3CCH J=20-19 and J=21-20 observations and
Gómez et al. (1991) derived the temperature of
90 K from
NH3(2, 2) and (3, 3) inversion transitions for the molecular envelope
surrounding the UC H II region.
We have observed the much lower CH3CCH transition J=5-4, which traces
the gas in outer layers and thus gives a lower rotational temperature
(34.6 K). The CH3CCH abundance estimate of
by Thompson &
Macdonald is an order of magnitude higher than our value of
.
NH3(3, 3), (4, 4), and (5, 5) observations by Acord et al. (1997) imply
a kinetic temperature of around 100 K for the outflowing gas and
they found a SiO abundance of
in this gas,
which is also an order of magnitude higher than the abundance estimate
in the present study in the position of W28 A2 (2). These results support
the idea that the CH3CCH and SiO abundances are clearly higher
in the hot core than in the warm envelope.
W31 is one of the largest H II complexes in our Galaxy.
G10.6-0.4 lying in the centre of our SIMBA map is one of the main extended
H II regions in this complex along with G10.2-0.3 and G10.3-0.1
(see, e.g., Corbel & Eikenberry 2004, and references therein).
NH3(1, 1) and (3, 3) observations show that the molecular core
in G10.6-0.4 is rotating and collapsing inward toward the UC H II
region where molecular masers are seen distributed in the following way:
H2O and CH3OH masers linearly in the plane of rotation, and OH masers
along the rotation axis (see, e.g., Sollins & Ho 2005 and references therein).
Further evidence for collapse comes from the two-peaked profile of the CO(4-3)
line observed with the APEX (Wyrowski et al. 2006).
Observations made with the VLA by Sollins & Ho (2005) suggest that there is
a rotating molecular accretion flow which flattens but does not form
an accretion disk.
Velocities of molecular lines observed from the surrounding neutral gas in
G10.6-0.4 suggest that accretion flow passes through the H II region
boundary and continues inward as an ionized flow (Keto 2002).
The associated infrared source (see Fig. 1),
IRAS 18075-1956, has a luminosity of
(Casoli et al. 1986).
The distance of W31 is not well known. However, spectral types of O stars identified in W31 and the NIR photometry performed by Blum et al. (2001)
show that the distance to W31 must be 3.7 kpc. Corbel and Eikenberry
(2004) suggest that G10.6-0.4 is located on the -30 km s-1 spiral arm
at a distance from the Sun of
kpc, which is consistent with
the spectrophotometric distance obtained by Blum et al.
We have used the value of 3.1 kpc which was derived by Blum et al. assuming
that O star spectra are consistent with the ZAMS. Note that
Sollins & Ho (2005) and Sollins et al. (2005) adopt the kinematic distance
of 6.0 kpc from Downes et al. (1980).
The W33 complex is heavily obscured in the visual. It contains a compact radio core G12.8-0.2. The FIR sources W33 A (G12.91-0.26) and W33 B (G12.70-0.17) associated with OH/H2O masers lie on opposite sides of G12.8-0.2 (Capps et al. 1978; Stier et al. 1984).
IR and radio continuum observations have revealed a cluster of compact sources
in the region of G12.8-0.2 (Dyck & Simon 1977; Haschick & Ho 1983).
The total IR luminosity of G12.81-0.19 assuming a distance of 3.7 kpc is
(Stier et al. 1984). According to the results
of Stier et al. the minimum mass of the star associated with
the radio core is 20
.
G12.8-0.2 is considered as a good
candidate for a very young OB cluster (see, e.g., Beck et al. 1998).
Spectral line observations at 3 and 2 mm, and dust continuum
observations at 1.2 mm performed with the SEST were used to derive
physical characteristics and fractional SiO and CH3CCH abundances in
15 high-mass star-forming cores. The sample represents typical GMC
cores with the following average properties
K,
,
and
.
While most cores seem to be near the virial equilibrium,
for five of them the gravitational potential energy
is large compared with the kinetic energy.
The fractional SiO abundances determined from the velocity range with
detectable CH3CCH emission were found to be
.
CH3CCH abundances were found to be
.
Consistent with earlier results towards GMC cores
(e.g., Ungerechts et al. 1997; Bergin et al. 1997)
the CH3CCH abundance is found to be larger than in the dense cores
of dark clouds. It also shows a slight increase as
a function of
.
As this tendency does not agree with the predictions of gas-phase chemistry
models (e.g. Lee et al. 1996), a possible explanation is that warm
conditions lead to an intensified desorption of the precursor
molecules of CH3CCH from the icy mantles of dust grains. This suggestion
is supported by the chemistry model of Canosa et al. (1997).
The SiO abundances are midway between the upper limit from dark clouds
and PDRs (10-12) and the values derived for powerful shocks
(
10-8). The abundance seems to decrease with
rising temperature. This finding contradicts models where
SiO production is dominated by neutral reactions with activation energies
(Ziurys et al. 1989; Langer & Glassgold 1990), and with models where
desorption of Si-containing species from grain mantles is significant
for SiO production (MacKay 1996). As suggested by
the observational results of Fuente et al. (2005), warmer cores
represent more evolved objects where highly
energetic protostellar outlows releasing SiO into the gas phase are less
frequent; this together with rapid post-shock processing
have lead to a diminished SiO abundance.
Acknowledgements
We thank the referee for a careful and critical reading of the manuscript and for very helpful comments and suggestions. The work has been supported by the Academy of Finland through grants Nos. 1201269 and 1210518. We acknowledge the use of the SIMBAD database, operated at CDS, Strasbourg, France, and the NASA's SkyView facility (http://skyview.gsfc.nasa.gov) located at NASA Goddard Space Flight Center.
Molecule | Transition | ![]() |
![]() |
29SiO | J=2-1 | 85 759.0000 | 6.195 |
28SiO | J=2-1 | 86 846.9600 | 6.252 |
29SiO | J=3-2 | 128 636.7064 | 12.390 |
28SiO | J=3-2 | 130 268.6100 | 12.504 |
CH3CCH | JK=50-40 | 85 457.2720 | 12.304 |
CH3CCH | JK=51-41 | 85 455.6220 | 19.505 |
CH3CCH | JK=52-42 | 85 450.7300 | 41.107 |
CH3CCH | JK=53-43 | 85 442.5280 | 77.112 |
CH3CCH | JK=60-50 | 102 547.9842 | 17.225 |
CH3CCH | JK=61-51 | 102 546.0241 | 24.426 |
CH3CCH | JK=62-52 | 102 540.1447 | 46.029 |
CH3CCH | JK=63-53 | 102 530.3487 | 82.033 |
Source |
![]() |
![]() |
l, b | d |
![]() |
![]() |
notes |
[h:m:s] | [
![]() ![]() ![]() |
[
![]() |
[kpc] | [kpc] | [km s-1] | ||
IRAS 12326-6245 | 12:35:34.1 | -63:02:28 | 301.13-0.22 | 4.4a | 7.1 | -39.5 | OHA/H2OB/CH3OHC/UC H IID,E |
G326.64+0.61 | 15:44:32.3 | -54:05:54 | 326.64+0.61 | 2.7 | 6.4 | -39.8 | H2OF |
OH328.81+0.63 | 15:55:47.3 | -52:43:08 | 328.81+0.63 | 2.9 | 6.2 | -41.7 | OHF/CH3OHG/UC H IIH |
IRAS 15566-5304 | 16:00:30.7 | -53:12:34 | 329.03-0.20 | 3.0 | 6.2 | -43.6 | OHI/H2OJ/CH3OHG |
G330.95-0.19 | 16:09:52.2 | -51:55:20 | 330.95-0.19 | 6.1 | 4.3 | -100.0 | OHH/H2OF |
G345.01+1.8N | 16:56:47.2 | -40:14:09 | 345.01+1.80 | 1.7 | 6.9 | -13.8 | OHK/H2OF/CH3OHG/UC H IIH |
IRAS 16562-3959 | 16:59:41.9 | -40:03:42 | 345.49+1.47 | 1.6 | 7.0 | -12.2 | OHH/UC H IID |
G345.00-0.23 | 17:05:11.7 | -41:29:11 | 345.00-0.23 | 2.9 | 5.7 | -27.7 | OHK/H2OF/CH3OHG |
NGC 6334 FIR-V | 17:19:55.9 | -35:57:45 | 351.16+0.70 | 1.7b | 7.2 | -6.6 | OHK/H2OF |
NGC 6334F | 17:20:53.5 | -35:47:01 | 351.42+0.65 | 1.7b | 7.1 | -7.0 | OHL/H2OF/CH3OHG/UC H II H |
G351.77-0.54 | 17:26:42.6 | -36:09:17 | 351.77-0.54 | 0.7 | 7.8 | -3.2 | OHK/H2OF/CH3OHG/UC H IIH |
G353.41-0.36 | 17:30:26.0 | -34:41:57 | 353.41-0.36 | 3.5 | 5.1 | -16.8 | OHM,K/H2OF/CH3OHF/UC H IIH |
W28 A2(2) | 18:00:30.4 | -24:04:00 | 5.88-0.39 | 2.7 | 2.4 | 6.1 | OHH/H2ON/UC H IIH |
W31 (2) | 18:10:28.7 | -19:55:50 | 10.62-0.38 | 3.1c | 5.5 | -2.8 | OHH/H2ON/UC H IIH |
W33 CONT | 18:14:13.6 | -17:55:25 | 12.81-0.20 | 3.7 | 4.9 | 34.8 | H2ON |
References: A Cohen et al. (1988); B Caswell et al. (1989); C Caswell et al. (1995); D Bronfman et al. (1996); E Zinchenko et al. (1995); F The catalog of non-stellar H2O/OH masers (Braz & Epchtein 1983); G Caswell (2000); H Caswell (1998); I Caswell et al. (1995); J Scalise et al. (1989); K Caswell & Haynes (1983); L Moran & Rodriquez (1980); Gaume & Mutel (1987); M Caswell et al. (1981); N Arcetri Atlas of galactic H2O masers (Comoretto et al. 1990; Palagi et al. 1993; Brand et al. 1994). Distance references: a Zinchenko et al. (1995); Osterloh et al. (1997); b Neckel (1978); c Blum et al. (2001). When no reference is given, the distance is a kinematic distance calculated from the CH3CCH velocities. |
Source | peak No. |
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IRAS 12326-6245 | 12:35:34.7 | -63:02:40 | 546 | |
G326.64+0.61 | 1 | 15:44:43.2 | -54:05:46 | 226 |
2 | 15:44:33.2 | -54:05:22 | 201 | |
3 | 15:44:59.5 | -54:02:26 | 197 | |
4 | 15:44:56.9 | -54:07:14 | 101 | |
5 | 15:45:01.4 | -54:09:14 | 63 | |
OH328.81+0.63 | 15:55:48.6 | -52:43:04 | 419 | |
IRAS 15566-5304 | 16:00:31.1 | -53:12:38 | 166 | |
G330.95-0.19 | 1 | 16:09:53.1 | -51:54:52 | 995 |
2 | 16:10:17.3 | -51:58:44 | 36 | |
G345.01+1.8N | 1 | 16:56:47.2 | -40:14:21 | 303 |
2 | 16:56:40.9 | -40:13:17 | 151 | |
3 | 16:56:43.7 | -40:15:57 | 65 | |
4 | 16:56:52.1 | -40:17:01 | 44 | |
IRAS 16562-3959 | 1 | 16:59:41.7 | -40:03:34 | 401 |
2 | 16:59:28.5 | -40:10:06 | 131 | |
3 | 16:59:06.9 | -40:05:50 | 49 | |
G345.00-0.23 | 17:05:11.0 | -41:28:59 | 372 | |
NGC 6334 | 1 | 17:20:53.4 | -35:46:58 | 1150 |
2 | 17:20:55.4 | -35:45:06 | 850 | |
3 | 17:20:19.3 | -35:54:51 | 519 | |
4 | 17:19:57.5 | -35:57:47 | 447 | |
5 | 17:20:32.4 | -35:51:23 | 147 | |
6 | 17:20:42.9 | -35:49:14 | 81 | |
7 | 17:20:15.3 | -35:59:39 | 61 | |
8 | 17:20:33.0 | -35:46:51 | 59 | |
G351.77-0.54 | 1 | 17:26:42.9 | -36:09:17 | 1007 |
2 | 17:26:39.0 | -36:08:05 | 114 | |
3 | 17:26:47.6 | -36:12:05 | 82 | |
4 | 17:26:25.8 | -36:04:45 | 55 | |
G353.41-0.36 | 17:30:26.7 | -34:41:45 | 487 | |
W28 A2(2) | 1 | 18:00:30.4 | -24:04:00 | 626 |
2 | 18:00:40.9 | -24:04:16 | 205 | |
3 | 18:00:15.2 | -24:01:20 | 40 | |
W31 (2) | 1 | 18:10:28.7 | -19:55:50 | 690 |
2 | 18:10:18.5 | -19:54:22 | 60 | |
W33 CONT | 1 | 18:14:13.9 | -17:55:41 | 1071 |
2 | 18:14:39.7 | -17:52:05 | 172 | |
3 | 18:13:54.8 | -18:01:49 | 119 | |
4 | 18:14:36.3 | -17:55:01 | 57 | |
5 | 18:14:25.7 | -17:53:57 | 53 | |
6 | 18:14:07.7 | -18:00:37 | 48 |
Source |
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K=0 | K=1 | K=2 | K=3 | |||
IRAS 12326-6245 | -39.54(0.04) | 3.68(0.02) | 0.69(0.02) | 0.58(0.02) | 0.29(0.02) | 0.13(0.02) |
G326.64+0.61 | -39.76(0.02) | 3.35(0.02) | 1.41(0.02) | 1.10(0.02) | 0.47(0.02) | 0.19(0.02) |
OH328.81+0.63 | -41.72(0.01) | 3.07(0.02) | 1.89(0.02) | 1.57(0.02) | 0.63(0.01) | 0.30(0.01) |
IRAS 15566-5304 | -43.64(0.03) | 3.73(0.03) | 0.95(0.03) | 0.77(0.03) | 0.31(0.03) | 0.11(0.02) |
G330.95-0.19 | -90.95(0.05) | 5.42(0.06) | 1.31(0.02) | 1.00(0.02) | 0.54(0.02) | 0.21(0.02) |
G345.01+1.8N | -13.78(0.02) | 3.60(0.02) | 1.23(0.02) | 0.96(0.02) | 0.48(0.01) | 0.24(0.01) |
IRAS 16562-3959 | -12.17(0.01) | 3.65(0.01) | 4.92(0.02) | 4.16(0.02) | 1.92(0.02) | 1.06(0.001) |
G345.00-0.23 | -27.69(0.04) | 5.56(0.05) | 1.53(0.02) | 1.36(0.02) | 0.46(0.02) | 0.26(0.02) |
NGC 6334 FIR-V | -6.63(0.01) | 3.45(0.01) | 2.08(0.02) | 1.61(0.02) | 0.76(0.02) | 0.39(0.02) |
NGC 6334F | -7.00(0.03) | 5.08(0.03) | 2.17(0.02) | 1.65(0.02) | 0.83(0.02) | 0.45(0.02) |
G351.77-0.54 | -3.16(0.02) | 4.88(0.02) | 3.10(0.03) | 2.56(0.02) | 1.13(0.02) | 0.62(0.02) |
G353.41-0.36 | -16.78(0.02) | 4.02(0.02) | 1.89(0.02) | 1.57(0.02) | 0.58(0.02) | 0.19(0.02) |
W28 A2(2) | 8.88(0.01) | 3.65(0.01) | 3.73(0.02) | 3.21(0.02) | 1.51(0.02) | 0.76(0.02) |
W31 (2) | -2.78(0.03) | 5.00(0.03) | 2.08(0.02) | 1.76(0.02) | 0.81(0.02) | 0.43(0.02) |
W33 CONT | 34.80(0.01) | 3.79(0.01) | 4.20(0.02) | 3.77(0.02) | 1.77(0.02) | 1.05(0.02) |
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Figure 2: The v=0, J=3-2 and v=0, J=2-1 lines of 28SiO and 29SiO towards each observed source. The weaker 29SiO spectra are multiplied by 2. |