Issue |
A&A
Volume 513, April 2010
|
|
---|---|---|
Article Number | A9 | |
Number of page(s) | 13 | |
Section | Interstellar and circumstellar matter | |
DOI | https://doi.org/10.1051/0004-6361/200912241 | |
Published online | 14 April 2010 |
A millimeter survey of ultra-compact
HII-regions and associated molecular clouds![[*]](/icons/foot_motif.png)
E. Churchwell1 - A. Sievers2 - C. Thum3
1 - Department of Astronomy, University of Wisconsin-Madison, 475 N.
Charter street, Madison, WI 53706, USA
2 - Instituto de Radio Astronomía Milimétrica, Avenida Divina Pastora,
7 Núcleo Central, 18012 Granada, Spain
3 - Institut de Radio Astronomie Millimétrique, Domaine Universitaire
de Grenoble, 300 rue de la Piscine, 38406 Saint Martin d'Hères, France
Received 1 April 2009 / Accepted 18 November 2009
Abstract
We report observations, using the IRAM 30 m telescope,
of 30 ultracompact and hypercompact
HII regions in the lines of HCO+(3-2)
and/or HCO+(1-0) and H
and/or H
.
Images are presented in both HCO+(3-2) and H
toward a subset of regions (16 in HCO+(3-2),
14 in H
)
with a resolution of 12''. In addition, H13CO+(3-2)
observations are reported toward 13 HII regions where
HCO+(3-2) displays complex profiles.
It is shown that the absorption dips in the HCO+ profiles
are due to HCO+ self-absorption, not absorption
of the HII free-free emission or warm dust emission
surrounding the HII region or two velocity components along
the line of sight. It was found that among the sources with
self-absorbed profiles, 8 are contracting and 5 are
expanding. Mass fluxes are found to be typically a few times 10-3
yr-1,
implying
time scales for massive star formation <105 yrs.
HCO+ and H2 column
densities are estimated for a subset of the sources from which masses
of the dense central cloud cores were
estimated. Implications of the derived column densities, masses, flow
velocities, and mass fluxes are discussed.
Key words: surveys - stars: formation - HII regions
1 Introduction
Ultra-compact (UC) HII regions are sites of recent massive star formation. UC HII regions are of special interest because they occupy an important stage in the evolution of young massive stars that is still only understood in broad outline. Identifying the ionizing star(s) and their associated natal cluster of lower mass stars in UC HII regions has proven quite difficult; identifications have been reported for only a few UC HII regions. The immediate regions around the central star(s) of UC HII regions are expected to be very dynamic due to possible infall, outflows, stellar winds, accretion disk rotation, turbulence, and shocks. The dynamics and physical properties of the photo-dissociation regions (PDRs) that surround UC HII regions and the ambient natal molecular gas that surrounds the PDRs must also be included to place UC HII regions in context with their environments. Spectroscopy of the H+, PDR(H0), and molecular envelopes (H2) is essential to understand the interactions of UC HII regions with their environments. Churchwell (2002) has argued, based on the lack of evidence for outflows in high resolution radio continuum images, that by the time young massive stars have formed UC HII regions they have largely ceased the accretion process. This is indirect evidence that has so far not been confirmed by spectroscopic observations. To understand the early evolution of massive stars and how they impact their environments, we must determine at what evolutionary stage young massive stars quench accretion and by what mechanism. Is it because the stars use all the matter in their neighborhood or the central radiation pressure over-powers gravity, or some other reason?
We also need to better understand the extent, morphology, and physical properties of the dense molecular and ionized gas in the immediate neighborhood of UC HII regions. Several papers have been published that address one or more of these issues toward specific UC HII regions (some of which are Baudry et al. (1981); Bourke et al. (1997); Cesaroni et al. (1991, 1994, 1998); Choi et al. (1993); Fey et al. (1992); Garay et al. (2007); Sandell & Sievers (2004); Zhu et al. (2008) and references therein), but a sensitive survey is not available that specifically probes both the dense molecular and ionized gas toward a substantial sample of UC HII regions with the same spatial resolution using a single-dish millimeter-wave telescope (i.e. sensitive to all scale sizes within the mapped area).
Here, we present observations of HCO+(3-2)
and/or (1-0) and H
and/or H
lines
toward 29 UC HII regions. The H
and/or HCO+(1-0) observations in the
3 mm band were made only toward the central positions of
7 UC HII regions. The HCO+(3-2)
and H
lines
in the 1 mm band were mapped toward 23 UC
HII regions (see Table 1). The
1 mm observations, we believe, represent a large enough sample
to begin to address systematics of the relationship between the ionized
gas and the surrounding molecular gas. Among these are:
is there systematic evidence for infall of molecular gas
toward UC HII regions? Is there evidence for
excessive line widths in the molecular or ionized gas and,
if so, what is the likely reason? Based on kinematic distance
determinations, is there evidence for a temperature gradient with
galactocentric radius? That is, are UC HII region
temperatures consistent with the observed metallicity gradients
(Churchwell & Walmsley 1975;
Shaver et al. 1983;
Maciel & Köppen 1994;
Mezger et al. 1979;
Afflerbach et al. 1996;
Afflerbach et al. 1997);
and, are there any systematics associated with the column densities of
HCO+ toward UC HII regions?
In the following we discuss the observations in Sect. 2, the data in Sect. 3, results in Sect. 4, and summary and conclusions in Sect. 5.
Table 1: Target sources and type of observation made.
2 Observations
Table 2: Spectral resolution and bandwidth available with the receiver/backend combinations used.
The bulk of the observations were made during the period of
28 March to 02 April 2006 using the IRAM
30 m telescope located at an altitude of 2850 m near
Granada (Spain). During most of this period the weather was good enough
for observations at 1.3 mm wavelength (
). We used the multibeam
1.3 mm receiver HERA (Schuster et al. 2004) whose two
orthogonal linearly polarized arrays were tuned to HCO+
at 267.557633 GHz and H
at 231.900942 GHz. Each array provides 3
3 pixels arranged as a center-filled square. The two arrays
are aligned to within
,
thus permitting precise relative pointing between the two transitions.
Each pixel consists of a diffraction-limited beam of 12'' (FWHP),
and the pixels are separated by 24''. The observations were
made in wobbler-switching mode where the reference positions are offset
in azimuth by 30'' to 120''. The array was kept aligned with
the equatorial system with the help of a derotator optical assembly
that compensates the field rotation due to changes of the parallactic
and Nasmyth angles. Maps sampled at 1/2 FWHP
intervals of size 66''
66'' size were obtained by stepping the telescope in
and
by 6, 12, and 18''. The 18 signals
generated by HERA were connected to three sets of backends. Their
spectral resolutions and bandwidths are listed
in Table 2.
During a small fraction of our observing time when the weather
was not good enough for 1.3 mm observations, we used the
Observatory's 3 mm single pixel receivers A100
and B100 tuned to H
at 106.737363 GHz and HCO+(1-0) at
89.188526 GHz. respectively. The VESPA correlator set to
spectral resolution of 78 kHz (0.26 km s-1
near 100 GHz) and bandwidth
of 140 MHz (480 km s-1
near 100 GHz) was connected to both receivers.
During a second session with the 30 m telescope on 03 August 2008 we observed the H13CO+(3-2) line toward a subset of our target list where the main isotope was found to be very optically thick. HC18O+(3-2) was observed in the few cases where even H13CO+(3-2) did not have a clean gaussian profile. We used the Observatory's single pixel receivers A 230 and B 230 tuned to the line frequencis of 260.25548 (H13CO+(3-2)) and 255.479389 (HC18O+(3-2)) GHz. As representative examples of a weak and a strong source, we show the spectra of the 3 HCO+ isotopes obtained toward G28.20-0.04N and G10.62-0.38 in Fig. 1. Note the presence of a weak SO2 emission line redward of HCO+ in G10.68-0.38. The spectrometers and their spectral characteristics are listed in Table 2.
The telescope beam is 12'' (full width at half power) at
1.3 mm. At the 3 mm transitions it is 26'' at
89 GHz and 24'' at 106 GHz. Pointing
observations were made every 2-3 h on nearby quasars; pointing
errors were found to be .
Line and continuum measurements are in units of antenna tempertaure,
,
calibrated in the usual hot/cold load technique with an estimated
precision of 10%. We use
to flux density conversion factors of 6.2, 8.6, and
10 Jy/K at 106, 232, and 268 GHz.
Our target sources were taken from various continuum surveys of ultra-compact HII-regions (Wood & Churchwell 1989; Shepherd & Churchwell 1996; Kurtz et al. 1994). We selected those sources which we estimated to have a 1.3 mm continuum flux density in our 12'' beam of at least 100 mJy. We extrapolated the measured 2 cm flux density to 1.3 mm assuming optically thin emission. Any optically thick emission, unrecognized at cm wavelengths, or a significant extended emission, not imaged by the VLA, of a halo surrounding the UCHII region, would increase the expected 1.3 mm flux density.
The observed sources are listed in Table 1. The last
column indicates the type of observations made: either HERA maps of the
HCO+ and H transitions or
pointed observations with the 3 mm single pixel receivers of
the transitions H
and HCO+(1-0). In the case of maps, the given
positions refer to map centers.
![]() |
Figure 1:
Spectra of the |
Open with DEXTER |
3 Data
3.1 Ionized gas
We detected recombination lines toward the peaks of 23 out of
30 (77%) of our target sources. Table 3 lists the line
parameters derived from Gaussian fits. The line widths are the full
widths at half power. The H lines are corrected
for instrumental broadening of 3 km s-1.
The H
lines
which were observed with 0.4 km s-1
spectral resolution do not need any such correction. For the sources
where no line was detected we give a
upper limit for a
spectral resolution of 2.6 km s-1.
The
km s-1
bandwidth used covers the full velocity range of galactic emission.
In most cases, the H brightness
distribution is only slightly broadened with respect to the beam.
As an example, Fig. 2 shows the UC
HII-region G10.62-0.38. Maps of the other regions are collected in
Fig. 3.
Linear baselines have been removed from all spectra in the map,
subtracting any free-free emission of the source. Note that the
coverage of these maps is slightly incomplete to the east due to one
bad pixel.
Table 3: Recombination line observations.
The H
line parameters were mostly derived from the HERA raster maps. The four
positions nearest to the nominal map center have offsets
and were averaged for the Gaussian fits, corresponding to a beam
smoothed to 14''. The H
parameters refer to
a 24'' beam.
![]() |
Figure 2:
Maps of G10.62-0.38 obtained with HERA at an angular resolution
of 12'' (circles at lower left corners) in the transitions of H |
Open with DEXTER |
In general, the H
line is only mildly affected by confusion with molecular lines. The
nearest known molecular transitions, an unidentified line and the
132-122 transition
of CH3C15N,
are -17 and +23 MHz, respectively, away from the
recombination line frequency. It is therefore only the line
wings which may be affected in a few particularly line-rich sources.
These cases are noted in Table 3. The kinematic
distances, given in Col. 6 of this table, are based on the H
LSR velocities assuming an orbital velocity of
220 km s-1. We adopt the IAU
standard distance from the Galactic center of 8.5 kpc,
even if this distance is somewhat larger than that currently
inferred from proper motions (Ghez et al. 2008).
![]() |
Figure 3:
H |
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![]() |
Figure 4: HCO+(3-2) maps obtained with HERA at an angular resolution of 12''. All maps in one row have been plotted using the same contour levels and color scale shown on the right. Contours are in K km s-1, start at and are in steps of 30 (row 1), 15 (row 2), 10 (row 3), 8 (row 4), and 4 (row 5). |
Open with DEXTER |
Table 4: HCO+ observations. Isotopes are abbreviated as 12 (HCO+), 13 (H13CO+), and 18 (HC18O+).
3.2 Molecular gas
The molecular clouds toward our target UC HII regions were
detected toward 24 of 28 sources (86%) observed via
their HCO+ emission. We failed to detect HCO+
within 30'' of 5 UC HII regions in our
sample (see Table 4).
In most cases, the HCO+(3-2) emission
is clearly resolved, its extent is in all cases larger than
that of H.
As an example, Fig. 2 shows the UC
HII region G10.62-0.38. Maps of the other regions are
collected in Fig. 4.
Table 4
lists the main properties of the HCO+ line:
the antenna temperature and radial velocity of the line peak,
the velocity of the self-absorption feature if there is any, and the
line width at 10% of the line peak. This quantity, while
little affected by self-absorption and still well
measurable, gives an indication of the presence of line broadening
mechanisms, like supersonic turbulence or high optical depth. In the
presence of self-absorption, the peak velocity is the measured H13CO+(3-2)
velocity. The measurement errors are typically K in
and
km s-1
for the weaker sources. We also give in Table 4 an estimate of the
size of the HCO+(3-2) core wherever reliable
data were obtained. Elliptical gaussians were fitted to the
velocity-integrated line intensity. The statistical errors of the major
and minor axis are 1'' or less, and less
than 5
for the position angle.
As in the case of the recombination lines, the HCO+(3-2) line parameters are derived from the four spectra of the raster maps closest to the nominal map center, and are thus derived for an effective beam of 14''. The HCO+(1-0) parameters refer to a 24'' beam (FWHP).
4 Results
A key goal in the quest to identify protostars that are actually in the
process of accreting mass is to find unequivocal evidence for mass
infall toward the center of molecular clouds. The period of active
accretion is believed to be short, typically
yr
for massive protostars. Consequently, we expect detection of a massive
protostar in the act of rapid accretion to be rare. Several surveys
have been undertaken to search for infall of molecular gas associated
with massive star formation. Among these are Fuller et al. (2005), Purcell
et al. (2006),
Klassen & Wilson (2007),
and Wu et al. (2007).
All of these used HCO+ and H13CO+(3-2)
observations along with a few other molecular probes. The search for
infall, for the most part relied on line profile properties specified
in the Myers et al. (1996)
``two-layer'' model and the sample was mostly toward molecular cores in
cold dark clouds with no bright continuum source such as UC
HII regions. Also, none of the surveys toward massive star
formation regions included simultaneous observations of a radio
recombination line (RRL) at the same spatial resolution,
so the velocities and spatial extent of the molecular gas
could not be compared with those of the associated HII region.
4.1 HCO+ and H30
relative velocities
In this section, we examine the motions of HCO+
relative to H.
The rest frequencies of the lines involved in this comparison are known
to a precision of at least 10 kHz, sufficient for deriving
velocities to a precision of
km s-1
as needed here. We assume that the velocity of H
represents the systemic velocity of the UC HII region. The HCO+
emission distributions are well
correlated with positions of the HII regions
(see Figs. 2-4). The HCO+
emission is generally more extended than that of the H
distribution.
Also, HCO+ and H
velocities are
similar, within a few km s-1
(Fig. 5).
This suggests that the HII regions and HCO+
gas are dynamically connected and that the HII regions are
probably embedded in the HCO+ clouds; we assume
that this is the case for all the sources in our sample where we have
both HCO+ and H
images.
Here, we will examine the relative velocities of HCO+
and H emission
to try to determine if the HII regions and molecular clouds
are separating from each other, or are at rest with respect to
each other. Interpretation of the kinematics from the HCO+
and HII velocities is complicated because it depends on how
optically thick the HCO+ line is
(i.e. we are mostly seeing only the nearside of the HCO+ cloud)
and whether the HCO+ profile has an
absorption dip.
Let us now consider the possible scenarios and the
implications for expected HCO+ line
profiles. If the HCO+ line is optically
thin, then we see motions from the entire line of sight.
In this case, if =
,
the bulk motions of the HII region and molecular
cloud are at rest relative to each other. If
,
then the molecular cloud and the HII region are in motion
relative to each other, the sense of which depends on the sign of
and v(RRL). However,
alone
cannot tell us if the molecular cloud is contracting
or expanding.
As we show in Sect. 4.2 below, the UC HII regions at 1 mm are very optically thin and are too faint to produce detectable absorption in the HCO+ line. This is also supported by the fact that absorption dips in the HCO+ profiles are detected more than a full half-power beam width away from the UC HII positions. Self-absorption, of course, requires that the HCO+ line be optically thick and that its excitation temperature decrease outward.
In the HCO+ optically thick scenario,
we see mostly the front face of the HCO+ cloud.
So a velocity difference between the H
and HCO+ lines could imply a relative
velocity between the HII region and the HCO+ cloud,
or contraction or expansion of the outer HCO+ gas
relative to the HII region. That is, we cannot
distinguish between relative bulk motions and contraction or expansion
of HCO+ about the HII region from
comparison of central line velocities. This can be resolved, however,
from line profile analysis using techniques outlined by Myers
et al. (1996)
and generalized by De Vries & Myers (2005).
Figure 5
shows the observed velocity difference between the central velocity of
the HCO+ cloud and that of the ionized
region for all sources where both velocities were reliably
observed. Whenever the HCO+ (3-2
or 1-0) line profile was distorted by an absorption
dip within the HCO+ emission profile,
the HCO+ central velocity was
obtained from H13CO+ profiles.
Of the 22 sources toward which both H13CO+(3-2)
and H
or both HCO+(1-0) and H
were detected, the HII region and HCO+ gas
are moving apart at speeds ranging from
km s-1
up to 9 km s-1; two sources
have relative speeds <0.5 km s-1.
Thus,
%
of the sample has significant relative motion between the
HII region and observed HCO+ gas.
The observed distribution of the velocity differences
may be compared to that expected in a scenario where the ionized gas is
streaming away from the edge of a molecular cloud. Such a scenario was
discussed for the Orion Nebula (Zuckeman 1973) where the
ionized gas happens to stream toward the observer at
km s-1.
Inasmuch as (i) the streaming directions in our
sample sources are random; (ii)
is the same for all sources; and (iii)
is a good representation of
the velocity of the molecular material,
we expect a distribution of velocity differences, projected on
the line-of-sight,
that
is flat between
.
Given the uncertainties due to the small size of our sample, this
scenario may well be what we see in Fig. 5. Nevertheless,
a small bias toward positive velocity differences is evident
from the figure. A departure from this simple scenario, which
predicts as many sources with
km s-1
as with
km s-1,
is further supported by the observed imbalance of
14 vs. 8 sources in these two velocity brackets.
Both of these trends may suggest an alternative scenario where
the UCHII regions are deeply embedded in their molecular
clouds. This geometry tends to reduce
since any streaming motions are
less asymmetric. If the HCO+ line
is (partially) optically thick, a small positive bias of
would result, as observed, if the molecular clouds
are mainly contracting. However, as noted above,
to determine if the HCO+ cloud
is contracting or expanding around the embedded HII region,
requires further analysis of the HCO+ profiles.
![]() |
Figure 5:
Histogram of the observed velocity differences between the ionized and
molecular gas, as obtained from our H |
Open with DEXTER |
4.2 Absorption dips in the HCO+ emission profiles
Of the 24 sources where HCO+ was detected, 14 have absorption dips within the emission profile, one of which, NGC 7538 IRS2, appears to have multiple absorption dips within the emission profile. Nine sources have no obvious absorption dip. Among those with observed HCO+ single absorption dips, the ``blue peak'' (i.e. lower velocity) of the double peaked HCO+ profile is brighter than the ``red peak'' toward 8 sources, and the red peak is brighter than the blue peak toward 5 sources (see Table 4). We refer to these profile types as B and R, respectively, in Table 4 and the rest of this paper.
Let us digress for a moment to summarize line profiles in an idealized collapsing molecular cloud. Myers et al. (1996) showed analytically that in the case of an optically thick, centrally condensed, collapsing cloud where the front and back halves have constant excitation temperature and dispersion velocities, the line profile will have two peaks separated by a self-absorption dip. For collapse, the blue peak will be brighter than the red peak (in our nomenclature a B-type profile). Since this analysis is symmetric, an absorption dip that falls in the blue half of the profile (an R-type profile in our nomenclature), would imply expansion. De Vries & Myers (2005) generalized the analytic model results using a Monte Carlo radiative transfer code that incorporated both a constant excitation temperature ``two-layer'' model and one in which excitation temperature increases inward as a function of optical depth. De Vries & Myers (2005) numerically calculated line profiles for collapsing clouds that fit a wide range of observed HCO+ profiles toward low-mass starless cores.
What is the origin of the absorption dips in the HCO+ lines? There are four possibilities: 1) absorption by HCO+ of the radio free-free continuum from the embedded HII region; 2) absorption by HCO+ of warm thermal dust continuum surrounding the HII region (heated by emission from the HII region); 3) self-absorption of HCO+ due to a negative temperature gradient with distance from the HII region; and, 4) two velocity components along the line of sight. Two velocity components along the line of sight can be ruled out by the single peaked H13CO+(3-2) profiles whose peak emission coincides with the velocity of the absorption dip in the H12CO+ profiles.
Let us now address whether the HCO+
absorption dips might be due to absorption of the HII region
free-free continuum. In principle, one could test this by
observing if the absorption dip is seen only toward the UC
HII region but disappears away from the HII region.
In this case, one could conclude that the absorption dip in
the HCO+ line is due to absorption of
the HII continuum emission. However, for
a 12'' HPBW at 267 GHz an
HII region of temperature 104 K
would have a brightness temperature of only (K)
3
10-9 EM (pc cm-6),
where EM is the emission measure. The typical emission measure
of UC HII regions are 107 to a few
times 108 pc cm-6
(Churchwell 1991).
Thus, the main beam brightness temperatures contributed by free-free
emission are a few tenths of a degree to perhaps
a degree, depending on the flux density of the source.
This is substantially less than that needed to produce the
observed absorption dips in the HCO+(3-2) lines.
This conclusion is strengthened by the fact that the
HII continuum distribution traced by the H
line
is a point source or only slightly extended in essentially all of our
sample, but absorption dips are detected a full HPBW
or more from the HII peak in all the sources in
our sample.
What about emission from warm dust surrounding the
HII region? Let us take a specific example, say
G10.62-0.38. The HCO+ profile shows
that at the minimum of the absorption dip,
is
about 6.5 K and the maximum
value
outside the absorption dip is
K. If the
HCO+ gas is moderately optically thick,
one would expect its excitation temperature to be approximately equal
to the kinetic temperature of the gas, implying
22 K. We will also assume that the maximum optical depth in
the HCO+ line
(all sources in Table 7 have HCO+(3-2) optical
depths ranging from 260 to 6 with an average value
of 77). We will assume that the warm dust shell around the
HII region is large enough that it fills the telescope main
beam (i.e.
). From the radiative transfer
equation, we solve for the continuum brightness temperature required to
produce the observed antenna temperature
at the absorption dip and find
![]() |
(1) |
Exploring the range of reasonable values of








where


Let us assume that the dust is moderately optically thick, say
and
the beam filling factor of warm dust is unity. Then solving
Eq. (2)
for
,
we find that
K
for the minimum
K.
The main result of this exercise is that for a reasonable range
of
(HCO+)
and
,
the dust temperatures would have to be unreasonably high to
produce the observed absorption dips. Although it is possible for dust
to achieve temperatures >1000 K within
HII regions near a hot star, it is quite unlikely
that dust could have such high temperatures a beamwidth or more from
the central HII region. We therefore conclude that the
absorption dips are most likely produced by HCO+
self-absorption due to high HCO+ optical
depths and a negative temperature gradient with distance from the
central heat source
(i.e. the HII region). The absorption dips are
spatially strongest toward the HII regions because this is
where the highest
(HCO+) values
are observed.
Two sources in our sample, SgrB2MC and G31.41+0.31, have absorption dips that go below the continuum level. This cannot be due to the source extending into our reference position, since our reference beam shows no evidence for residual emission toward either source. In the case of G31.41, our map of HCO+ shows that the cloud is small relative to the beam throw. We therefore must find another explanation for the subcontinuum absorption dips. Since molecular clouds are effectively bathed in an isotropic radiation field of 2.7 K, the absorption dips cannot go below 2.7 K. Also, scattering and absorption of line photons cannot remove more photons from the line than are produced in the line and therefore cannot account for absorption below the continuum. As we showed above, free-free emission from the HII regions can contribute a few tenths of a degree or so to the continuum and therefore could contribute an equivalent amount of absorption below the contnuuum. Also, as we showed above, thermal dust emission essentially contributes no measurable emission in the HCO+ line because of large line optical depths; however, outside the line frequencies dust could be the main contributor to the continuum emission. In the case where the dust has a negligible temperature gradient (such that it is not apparent in beam-switching mode), scattering and absorption of continuum photons (i.e. free-free and thermal dust) in the line could produce absorption dips below the continuum down to the minimum of 2.7 K if the line optical depths are large enough. Further support for scattering of both line photons and continuum photons at line center is the fact that the widths of the HCO+ lines are substantially broader than other molecular lines along the same line of sight that are presumably optically thin, including the H13CO+(3-2) line (compare the 12C and 13C isotopic line widths in Table 4). For treatments of radiative transfer in very optically thick lines see Auer (1968) and Auer & Mihalas (1972).
Table 5: Flow velocities, column densities, and mass fluxes
4.3 Infall or outflow mass fluxes
In Fig. 1,
we show the spectra of the three HCO+(3-2) isotopes
toward two UCHII regions as examples of both a strong and a
weak source where HCO+ is strongly
self-absorbed. The close
correspondence between the velocities of the absorption dips in the HCO+(3-2)
and the H13CO+(3-2) profiles
with the emission peak of the HC18O+(3-2) line
supports the conclusion that the double peaked profiles are due to
self-absorption and not to two velocity components along the line of
sight. Toward G28.20N, G81.67, and possibly G33.13, the H13CO+(3-2) profiles
appear to be weakly self-absorbed but weak enough that the profiles do
not appear to be strongly distorted. We use the H13CO+(3-2) profiles
to obtain the dispersion velocity of HCO+ and
use the H12CO+(3-2) profiles
to measure the parameters required to estimate the sign and magnitude
of the HCO+ mass flow around each
HII region using Eq. (9) in Myers et al. (1996). Results of our
analysis of self-absorbed HCO+ profiles
are reported in Table 5
where the measured H13CO+(3-2) line
full-width at half-maximum intensity,
,
is given in Col. 2. The line dispersion velocity in
Eq. (9) of Myers et al. (1996) is given in
Col. 3;
is
related to
by
(i.e.
is the half-width at 1/e intensity). The temperature of the
blue peak above the minimum of the absorption dip T(BD)
is given in Col. 4 of Table 5, the
temperature at the minimum of the absorption dip T(D)
is given in Col. 5, the temperature of the red peak T(RD)
is given in Col. 6, the velocity of the red peak
is given in Col. 7, the velocity of the blue peak
is given in Col. 8, the derived flow velocities
in Col. 9, the
half-power radius of the HCO+ emission
distribution in Col. 10, the H2 column
denisity determined from the H13CO+(3-2) emission
toward the HII region in Col. 11, and the mass fluxes
are
given in Col. 12.
The propagated uncertainties for
are based on an
% error
for
,
K
for
,
,
and
,
and
km s-1
for
and
.
The propagated uncertainties for
are based on the uncertainties for
,
%
for R, and 50% for N(H2).
The large uncertainties for N(H2)
are due to uncertainties in the conversion of NJ=3(H13CO+)
to
(H13CO+)
to N(H2); the
uncertainty of N(H2)
could be as much as an order of magnitude which would translate to a
similar uncertainty in
.
However, the very large optical depths in the HCO+ lines
(ranging from 260 to 6 with an average
of 77) require values of N(H2) >
1023 cm-2, which
is consistent with the values given in Tables 5 and 7 based on
the observed H13CO+(3-2) lines.
Thus it seems unlikely that we have over-estimated N(H2)
and
.
Negative values of flow velocities
imply outflow and positive values infall. Five sources apparently have
a mass outflow and eight infall. It is clear that the flow
velocities in these regions are generally larger than those found in
the lower-mass star formation regions studied by Myers et al. (1996) and others.
This is not surprising because massive star formation regions
are more dynamical due to much larger masses (gravity), radiation
fields, and wind luminosities. Presumably those having outflow profiles
are more evolved and have had time to reverse
infall to an outflow in five of the 13 sources that we are
able to analyze. Assuming spherical contraction or expansion, we use
the derived flow velocities (Col. 9, Table 5) to estimate
the mass flux in each region assuming spherical flow and constant flow
velocity with radius
![]() |
(3) |
where








Of the 8 sources that have self-absorbed HCO+ profiles, 5 have expansion motions. Exploration of the mechanism by which infall motions are reversed is beyond the scope of this paper. Presumably, these sources are more evolved and have had time to reverse core contraction via radiation pressure or stellar winds or both. It is not known what determines the final mass of massive stars, although much speculation has been given to this question. Detection of expanding young massive cores obviously has important implications for the termination of protostellar accretion and further study of this phase is needed.
Table 6: Column densities from single peaked HCO+ profiles.
Table 7: Column densities N and masses M from H13CO+ lines.
It is of interest to compare the flow direction of the
molecular gas with that of the ionized gas as discussed in
Sect. 4.1.
Comparing
with
(HCO+ - H
)
for each of the 22 sources where both lines are measured does
not show any significant trend. Both signs of
(HCO+ - H
)
are equally likely for the molecular cores where we see infall or
outflow motions. It seems that the large relative velocities
of the ionized gas (
km s-1),
directed in random directions away from the molecular cores dominates
the small infall or outflow motions (mostly
1 km s-1)
of the neutral material, thus masking any possible correlation.
4.4 HCO+ column densities
The column density toward optically thin sources in the upper
rotational state of a linear molecule is given by
NJu | = | ![]() |
|
![]() |
(4) |
where







![]() |
(5) |
using J(















![]() |
(6) |
for J(





Applying the appropriate relations to H13CO+(3-2) data
toward the peak of the HCO+ distribution,
we have estimated the total H2 column
densities in Table 5
(Col. 11) and Table 7
(Col. 8) for those sources with double peaked HCO+ profiles.
For an optically thick line, a lower limit for the
excitation temperature is
2.7 K +
/
;
this assumes that the line optical depth
is substantially larger than
unity (
). As
increases, the line optical
depth decreases, so there is a
lower limit on
set by the maximum line optical depth (i.e. at the minimum of
the absorption dip). As
increases above this lower
limit, the total HCO+ column
density will decrease due to dominance of the
term in the partition function. Using the expression for
above, the HCO+(3-2) lines all
have
well in excess of 3
as expected if the relationship between
and
above holds. As noted above, the sources in Table 7 have
optical depths in the HCO+(3-2) line
ranging from 6 to 260. The column densities of HCO+
and H2 may be somewhat over-estimated, but as we
have already noted there does not seem to be a significant additional
contribution to
from free-free emission although emission from extended cool dust may
make some contribution. It is unlikely that the molecular
column densities are over-estimated by as much as an order of magnitude
by an under-estimate of
.
Among the sources where HCO+ was
detected, 13 are self-absorbed in the HCO+(3-2) line
and very optically thick, so we had to use the H13CO+(3-2) line
to estimate column densities as described above. Six sources, however,
are not self-absorbed and have approximately Gaussian H12CO+(3-2) profiles.
They may still suffer some optical depth effects, but we can
approximate lower limits for column densities using Eq. (5).
Results for the 6 sources with no self-absorption dip and
approximate Gaussian profiles are given in Table 6 where the
measured peak H12CO+(3-2) line
temperature is given in Col. 2, the line full-width at half
maximum, FWHM, is in Col. 3, the HCO+ column
density in J=3 is in Col. 4, the total HCO+ column
density for
2.7 + 1.92
(K) is in Col. 5, and
the H2 column density assuming [H2]/[HCO+] =
(see e.g. Fuente et al. 2003) is in
Col. 6. For these sources, both N(HCO+)
and N(H2) are about an order
of magnitude lower than those with self-absorbed profiles,
as expected for more optically thin lines. However, we caution
that the values reported in
Table 6
are lower limits because the HCO+ lines
may not satisfy the optical thin assumption even though they do not
show evidence of self-absorption.
In Table 7,
we report the parameters used to derive total HCO+
and H2 column densities and core masses
from the observed H13CO+(3-2) line.
The same source size R is assumed
as for the main isotope (Table 5). We adopt the
kinematic distances given in Table 3 based on the H line
radial velocities, because the H
line is likely to be
more representative of the systematic velocity of the core than HCO+
which is probably affected by infall/outflow.
In Table 7,
Cols. 2 and 3 are the H13CO+(3-2) line
FWHM and peak line antenna temperature above the continuum;
Col. 4 is the H13CO+ column
density in the J=3 level; the lower limit
on the excitation temperature (
2.7 K + 1.92
)
is given in Col. 5; the upper limit on the total column
density of
H13CO+ is given in
Col. 6 and the total molecular hydrogen column density
assuming [H2]/[HCO+] =
is given in Col. 7. In Col. 8, an upper
limit on the core mass is given using M =
N(H2)/2R
where we assume for the mean molecular weight
and
=
1 amu. The spherical approximation seems reasonable from the
observed HCO+ emission distributions
(Figs. 2b
and 4).
The errors for N(H2) have
been discussed in Sect. 4.3
and dominate the estimated errors for the mass M.
We assign the lower limits of
%, but note that the errors
could be higher due to uncertainties in the conversions from N(H13CO+)
to N(H2).
As pointed out in Sect. 4.3, it is
more likely that both N(H2)
and M are under-estimated rather than
over-estimated.
With all the assumptions that have gone into the
determinations of total column densities and the masses determined from
the column densities, one should view the values as order of magnitude
estimates. It is likely that the masses are, in fact,
only lower limits on the total mass of the
natal H2 cloud since the HCO+ emission
probably only traces the dense central cores of more extended clouds.
The total mass of the natal clouds are of interest, because it is
important to determine the minimum cloud mass required to form a
massive O star with its associated cluster of lower mass
stars. Typically it has been found that
seem to be required to produce one intermediate mass O-star with its
accompanying cluster of lower mass stars (Churchwell 1997). The values
found here support thresholds of this magnitude when we take into
account that the HCO+ data probably
only sample the central cores of substantially larger molecular clouds.
G81.68+0.54 may be an exception to this.
4.5 H30
line widths
The H
line widths range from 21.5 to over 57 km s-1
in our sample of sources. In terms of the equivalent Doppler
temperature (i.e. the kinetic temperature that hydrogen gas
would have to have to produce the observed full-width-at-half maximum
(FWHM) linewidth, this range corresponds to 5000 K to
>35 000 K, respectively. The upper extreme of
this range is well outside the values typically found in
HII regions; temperatures as low as 5000 K are found
in the inner Galaxy where metallicities are substantially greater than
at the solar circle. 35 000 K is much hotter than the
kinetic temperatures that typical Galactic plane metallicities will
permit except in shocks and very near hot stars. A further
important point is that the H
line is unlikely to
suffer significant pressure broadening (see Gordon & Sorochenko
2002), even in
dense UC HII regions. Sewilo et al. (2008) has
successfully separated thermal, turbulent, pressure, and large-scale
motions in the hypercompact (HC) HII region G28.20-0.04N and
showed that this ultra-dense HII region cannot have
significant pressure broadening of the H
line and accommodate
the other broadening components that were independently measured. We
therefore will assume that H
lines with FWHM in
excess of
km s-1
is not due to pressure broadening, but require large scale motions such
as expansion or contraction, bipolar outflows, rotation (of an
accretion disk, torus or shell), and/or shocks. Turbulence, although a
contributor, was found by Sewilo et al. (2008) to be small
relative to large scale motions in G28.20-0.04N. It is no
surprise to find bulk motions in UC and HC HII regions, which
are expected to be very dynamic especially during the period of rapid
accretion that occurs as massive star are built. It is
therefore not difficult to understand, in principle, the large
linewidths found in our sample (see Sewilo et al. 2008, for a break
down of the various line broadening components
in G28.20-0.04N).
Electron temperatures of the HII regions could not be measured because the millimeter free-free continuum could not be reliably separated from the dust emission.
5 Summary and conclusions
We have presented observations of 30 ultracompact and hypercompact
HII regions in the lines of HCO+(3-2)
and/or HCO+(1-0) and H
and/or H
.
Images are presented in HCO+(3-2) and H
emission
regions and sizes are reported with a resolution of 12''.
In addition, H13CO+(3-2)
observations are reported toward 13 HII regions where
HCO+ profiles showed signs of
self-absorption. All data have been obtained with the IRAM
30 m telescope, mostly at 1 mm.
It was shown that the HII regions and HCO+
are generally in motion relative to each other at speeds ranging
from 0.5 to <9 km s-1.
Since the H
and HCO+ regions coincide almost
precisely, the relative velocities may represent a combination of
relative bulk motions between
the HII region and the molecular cloud plus
contraction/expansion of the HCO+ envelope.
We examined four possibilities for the origin of the
absorption dips in the HCO+ profiles:
HII free-free continuum, warm dust continuum surrounding the
HII region, two velocity components along the line of sight,
and HCO+ self-absorption.
It was shown that UC HII regions at 267 GHz
are too optically thin to produce bright enough continuum emission to
produce the observed absorption dips. For dust to produce a detectable
absorption dip it would have to have temperatures
K,
which we rejected because the absorption dips are seen more than one
full beam width away from the HII regions where it is highly
unlikely that dust could have such high temperatures. The hypothesis of
two velocity components along the line of sight was rejected based on H13CO+(3-2) profiles
which are not double peaked. Also, the H13CO+ line
peaks at the velocities of the absorption dips in the H12CO+ lines.
We therefore conclude that the absorption dips
are due to HCO+ self-absorption.
Lower limits on HCO+ column densities
were determined for sources that have gaussian profiles
(i.e. not double peaked) from the intensities of the HCO+ lines.
Column densities for sources with self-absorbed profiles were
determined using the observed H13CO+ lines;
to estimate H2 column
densities it was assumed that [HCO+]/[H13CO+] = 40
(see e.g. Langer & Penzias 1990) and [H2]/[HCO+] =
.
Masses for the dense cloud cores traced by HCO+ were estimated from the
observed diameter of the HCO+ emission
and the H2 column densities. Due to the
uncertain assumptions involved, the column densities and masses should
be considered order of magnitude estimates. Evenso, they are useful
because they support independently
determined threshold masses required to form massive stars along with
their associated lower mass stars.
Using the two-layer analytic model of Myers et al. (1996) applied to the
13 sources that have double peaked HCO+ profiles,
we derived mass flow velocities and mass flux rates. It was
found that 8 sources have infall velocities and 5
have outflow velocities; this is similar to
the fractions of expanding and contracting cores found in the Orion
molecular cloud by Velusamy et al. (2008). The flow
velocities are typically a few tenths of km s-1,
although two sources (G10.62 and W51D) have flow
velocities km s-1.
Expanding young massive protostellar cores have important implications
for the termination of protostellar accretion and the final mass of a
massive protostar. An obvious mechanism to reverse infall is
the action of stellar radiation and winds; however, this needs to be
explored more thoroughly to be put on a more solid theoretical and
observational foundation.
Mass flux rates were found to be quite large, typically a few
times 10-3 yr-1.
Both the flow velocities and the mass fluxes found here are consistent
with those found by Barnes et al. (2010) in the massive
protostellar cluster By 72 in Carina. At these rates,
a 40
star could be formed in <105 yr.
Such short time scales are consistent with massive star formation model
predictions and inferences from independent observations. Such short
time scales
imply that detection of protostars in the rapid accretion phase should
be rare. Presumably, the reason we have been successful in
detecting 8 such sources out of a sample of 30 is
because the sample has been carefully selected for especially young
massive star formation regions.
The observed radio recombination lines could not be used to derive electron temperatures of the HII regions because we were unable to reliably detect the free-free continuum emission near 267 GHz.
AcknowledgementsWe thank the IRAM Director, Pierre Cox, for granting additional telescope time which permitted us to conclude this project in a timely manner. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain). IRAM telescope staff efficiently supported the observations. We also thank an anonymous referee for a very careful reading of the mauscript and for several suggestions that have improved this paper. E.C. acknowledges partial support from NSF grant AST-0808119 and NASA contract No. 1282620.
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Footnotes
All Tables
Table 1: Target sources and type of observation made.
Table 2: Spectral resolution and bandwidth available with the receiver/backend combinations used.
Table 3: Recombination line observations.
Table 4: HCO+ observations. Isotopes are abbreviated as 12 (HCO+), 13 (H13CO+), and 18 (HC18O+).
Table 5: Flow velocities, column densities, and mass fluxes
Table 6: Column densities from single peaked HCO+ profiles.
Table 7: Column densities N and masses M from H13CO+ lines.
All Figures
![]() |
Figure 1:
Spectra of the |
Open with DEXTER | |
In the text |
![]() |
Figure 2:
Maps of G10.62-0.38 obtained with HERA at an angular resolution
of 12'' (circles at lower left corners) in the transitions of H |
Open with DEXTER | |
In the text |
![]() |
Figure 3:
H |
Open with DEXTER | |
In the text |
![]() |
Figure 4: HCO+(3-2) maps obtained with HERA at an angular resolution of 12''. All maps in one row have been plotted using the same contour levels and color scale shown on the right. Contours are in K km s-1, start at and are in steps of 30 (row 1), 15 (row 2), 10 (row 3), 8 (row 4), and 4 (row 5). |
Open with DEXTER | |
In the text |
![]() |
Figure 5:
Histogram of the observed velocity differences between the ionized and
molecular gas, as obtained from our H |
Open with DEXTER | |
In the text |
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