Outflow and Infall in a Sample of Massive Star Forming Regions.
P. D. Klaassen & C. D. Wilson
Dept. of Physics and Astronomy, McMaster University, Hamilton, ON, Canada
klaassp@physics.mcmaster.ca
ABSTRACT
We present single pointing observations of SiO, HCO+ and H13CO+ from
the James Clerk Maxwell Telescope towards 23 massive star forming regions
previously known to contain molecular outflows and ultracompact HII regions.
We detected SiO towards 14 sources and suggest that the non-detections in the
other nine sources could be due to those outflows being older and without ongoing
shocks to replenish the SiO. We serendipitously detected SO2 towards 17 sources
in the same tuning as HCO+. We detected HCO+ towards all sources, and
suggest that it is tracing infall in nine cases. For seven infall candidates, we
estimate mass infall rates between 1×10−2 and 2×10−5 M⊙ yr−1. Seven sources
show both SiO detections (young outflows) and HCO+ infall signatures. We also
find that the abundance of H13CO+ tends to increase along with the abundance of
SiO in sources for which we could determine abundances. We discuss these results
with respect to current theories of massive star formation via accretion. From
this survey, we suggest that perhaps both models of ionized accretion and halted
accretion may be important in describing the evolution of a massive protostar
(or protostars) beyond the formation of an HII region.
Subject headings: Stars: Formation — ISM: Jets and Outflows — Accretion —
HII regions — Submillimeter — Molecular Processes
1. Introduction
The dynamics in massive star forming regions are, in general, much more complex
than in regions which form only low mass stars. For instance, in the early stages within
a low mass star forming region, the dynamics can be understood in terms of a few broad
categories: large scale infall, which causes a disk to form, accretion through the disk, and
outflow to release angular momentum (see for example, Di Francesco et al. 2001, André
et al. 1993, Muzerolle et al. 2003). In intermediate and high mass star forming regions,
http://arxiv.org/abs/0704.1245v1
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turbulence, stellar winds, multiple sites of star formation, and, for regions with massive star
formation, the presence of HII regions, all contribute to the dynamics in these regions as
well (i.e. Beuther et al. 2006, Shepherd & Churchwell 1996, McKee & Tan 2003, Krumholz
et al. 2005). The more complicated source dynamics make the processes involved in the
formation of the most massive stars much more difficult to understand than those involved
in the formation of lower mass stars. Adding to the complexity, massive stars do not form as
often as their lower mass counterparts and so we must look to larger distances before finding
examples of high mass star formation. For instance, the average distance to the 63 sources
in Shirley et al. (2003, hereafter S03) is 5.3 kpc.
If we assume massive stars form through accretion, that this accretion occurs in the
inner regions of disks (i.e. Pudritz & Norman 1986, or Shu et al. 1994), and that these
disks have radii of a few thousand AU (Chini et al. 2004, Beltran et al. 2004, Cesaroni et
al. 2005), we do not yet quite have the resolving power to detect accretion directly at the
distances to massive star forming regions (1000 AU at 5.3 kpc is ∼ 0.2′′). Here, we define
accretion as the infall motions from the disk onto the forming star, in contrast to the larger
scale motions of envelope material falling onto the disk. However, while we cannot observe
accretion directly, its presence can be inferred from the presence of accretion tracers such
as larger scale infall and outflow. Infall can act to replenish disk material as mass accretes
onto a protostar (Nakamura 2000), while molecular outflows serve as a release mechanism for
the angular momentum which builds up during the accretion process (e.g Arce et al. 2006).
These large scale motions are seen in star forming regions of all mass scales (see for instance
Beuther & Shepherd 2005).
In massive star forming regions, the accretion rates are orders of magnitude greater than
in low mass star forming regions (i.e. Beuther et al. 2002), while the accreted masses are
only approximately one order of magnitude greater. These accretion rates and masses result
in accretion timescales that are much shorter than in low mass star forming regions, which
allows the Kelvin-Helmholtz timescale to become important in the evolution of the protostar
(e.g., there is no pre-main sequence stage for massive star formation). The outward radiation
and thermal pressure from the forming star becomes strong enough that it can ionize the
surrounding medium and a small, highly ionized HII region (either hypercompact (HCHII) or
ultracompact (UCHII) region, Keto 2003) can form. It is still unclear whether the outward
pressure needed to create the HII region is strong enough to halt accretion, or whether
accretion can continue in some form (either through a molecular or ionized disk, or through
an ionized accretion flow) after the formation of an HII region. Some models suggest that
accretion must halt before the onset of a visible UCHII region (i.e. Garay & Lizano 1999,
Yorke 2002), while other models suggest that an ionized accretion flow can continue through
an HII region (i.e. Keto 2003, 2006).
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There is now also observational evidence which suggests that, once the protostar be-
comes hot enough to ionize its surroundings, both modes of massive star formation (halted
and ionized accretion) are possible. G10.6-0.4 has been shown to have an ionized accretion
flow by Sollins et al. (2005) and Keto & Wood (2006), while accretion in G5.89-0.39 seems
to have halted at the onset of the UCHII region (Klaassen et al. 2006). Although sample
statistics at this point are still quite small, these two examples pose interesting questions.
We do not yet have enough data to determine whether the apparently conflicting models of
halted and ionized accretion can both be correct. However, we can begin with a uniform
survey of infall and outflow tracers in massive star forming regions in order to constrain
massive star formation scenarios.
In this paper, we present a survey of 23 massive star forming regions. Because we
are interested in the relationship between accretion and outflow after the formation of an
HII region, our source selection criteria include (1) the presence of an UCHII region, which
indicates that there is a massive protostar forming, and (2) previous evidence of outflows,
which suggests ongoing accretion in most formation scenarios. Sources were selected based
on inclusion in the Wood & Churchwell (1989) and Kurtz et al. (1994) catalogs of UCHII
regions as well as having molecular outflow signatures in the Plume et al. (1992) survey of
massive star forming regions. Additional sources were taken from Hunter (1997) which were
shown to have both UCHII regions and molecular outflows.
We describe the observations collected for this survey in Section 2, we discuss the results
of these observations in Section 3, and present our conclusions in Section 4.
2. Observations
Observations of SiO (J=8-7), HCO+, and H13CO+ (J=4-3) were taken at the James
Clerk Maxwell Telescope (JCMT)1 in 2005 (as parts of projects M05AC11 and M05BC04).
SiO (347.330 GHz) and H13CO+ (346.999 GHz) were observed simultaneously in the same
sideband by tuning the receiver to 347.165 GHz. Thirteen or twenty minute observations,
depending on the source elevation, were taken towards each source with a velocity resolution
of 1.08 km s−1, which resulted in rms noise levels of TMB <0.07 K. Separately, we observed
HCO+ (356.370 GHz) with a velocity resolution of 0.53 km s−1 to an rms noise limit of
TMB <0.13 K in twenty minute integrations. Both sets of observations were taken in position
1The James Clerk Maxwell Telescope is operated by The Joint Astronomy Centre on behalf of the Particle
Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organisation for Scientific
Research, and the National Research Council of Canada.
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switching mode with dual mixers and a sideband rejection filter in place. Table 1 shows the
positions, rms noise limits of both tunings, the local standard of rest velocity, and distances
to all sources in this survey. The half power beam width for these observations is 15′′, and
the main beam efficiency is ηmb=0.62. Data were obtained using the DAS autocorrelator
system and reduced using the SPECX software package.
Linear baselines were removed from all spectra except for those towards G10.47. For
this source, there were so many different chemical species in the observed spectrum that we
were unable to fit a linear baseline over the entire ∼ 700 km s−1 bandwidth of the 347 GHz
observations or the ∼ 450 km s−1 bandwidth of the 356 GHz observations. In this case, no
baseline was removed.
We also present 9 × 9 maps of one source (G45.07) in the same emission lines. These
raster maps are sampled every 5” and have rms noise limits of 0.10 and 0.14 K (TMB) for
the 347 and 356 GHz tunings, respectively. Note that these values are different than the
ones reported in Table 1 for the single pointing observations. The DAS autocorrelator was
configured with the same tunings as were described above for the single pointing observations,
and the maps were centered at the same position.
3. Results
This survey of single pointing observations towards 23 massive star forming regions is
meant as an initial, uniform survey from which to base future observations. With these
observations, we can only comment on the molecular gas component within our beams; we
cannot discuss the larger scale molecular dynamics, or the ionized gas components of these
regions. For our sources, the selection criteria of having an HII region confirms that these
regions are forming massive stars.
Figures 1 through 4 show the single pointing SiO and HCO+/H13CO+ spectra towards
all sources; the spectra are ordered according to SiO integrated intensity. For each panel in
these figures, line brightnesses have been corrected for the JCMT main beam efficiency and
centered on the VLSR of the source (Table 1). Figures 1 and 2 show SiO and HCO
+/H13CO+
spectra, respectively, towards the sources with no SiO detections. Figures 3 and 4 show the
SiO and HCO+/H13CO+ spectra, respectively, towards sources with SiO detections. Peak
line strengths and integrated intensities are given for all three lines in Table 2.
SiO was only detected in 14 out of our 23 sources, where we define a detection as a
minimum of 4 σ in integrated intensity. The rms noise limits in integrated intensity were
calculated using ∆I = Trms∆v
Nchan where Trms is the rms noise level in K, ∆v is the velocity
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resolution of the observations, and Nchan is the number of channels over which the integrated
intensity is calculated. HCO+ was detected in all sources and H13CO+ was detected in all but
two sources. Along with HCO+, we serendipitously observed SO2 (J=104,6-103,7 at 356.755
GHz) in 17 of our sources. The HCO+ and SO2 lines are only separated by 17 km s
−1 and
thus the lines were blended in eight sources. For the sources with SO2 detections, we have
also plotted (in gray) the lower intensity HCO+ observations in order to highlight the SO2
emission (Figures 2 and 4).
Double peaked HCO+ line profiles were observed towards 10 sources, with nine of them
having stronger blue peaks than red. This blue line asymmetry in an optically thick tracer
such as HCO+ is often suggestive of infall (i.e Myers et al. 1996). We discuss the possibility
of our observations tracing large scale infall further in Section 3.2.
The distances to our sources, as taken from the literature, are shown in Table 1. The
average distance is 5.7 ± 3.8 kpc, where the error quoted reflects the 1 σ dispersion in the
distances. Since our observations were taken with a 15′′ beam, this resolution corresponds
an average linear size of 0.4 pc for our observations.
SiO is a well known outflow tracer, since in the general interstellar medium, Si is frozen
out onto dust grains. When the gas in a region is shocked (i.e. the gas through which a
protostellar outflow is passing) the dust grains can sublimate and Si is released into the gas
phase. After the passage of a shock, the SiO abundance ([SiO]/[H2]) can jump to almost
10−6, whereas the dark cloud abundance of SiO is often closer to 10−12 (see for example
Schilke et al. 1997, Caselli et al. 1997, or van Dishoeck & Blake, 1998).
While SiO is easily identified as an outflow tracer, the emitting region for HCO+ is much
less certain. Many authors suggest that HCO+ can be used to trace the envelope material
surrounding a protostellar region (i.e. Hogerheijde et al. 1997, Rawlings et al. 2004), while
others suggest that it traces disk material (i.e. Dutrey et al. 1997). One thing that is
apparent, however, is that it becomes optically thick very quickly and readily self absorbs.
We detected SO2 in 17 of our sources, suggesting that our beam contains at least some
molecular gas at temperatures greater than 100 K (see for instance Doty et al. 2002, Charnley
1997). However, Fontani et al. (2002) have determined the average temperature in twelve
massive star forming regions to be 44 K, using observations with beam sizes comparable to
those presented here. For seven of our sources, which were observed in the Fontani et al
(2002) sample, the average temperature is also 44 K. Thus, in the following analysis, we
adopt an ambient temperature of 44 K for all sources.
Table 3 shows the column densities for each region derived for both SiO and H13CO+.
The column density was calculated assuming that each tracer is optically thin, in local
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thermodynamic equilibrium, and at an ambient temperature of 44 K. For optically thin lines,
the column density of the observed transition scales directly with the integrated intensity of
the line (see for instance, Tielens 2005):
8kπν2
TMBdv (1)
where Nu is the column density in the upper state of the transition, Auℓ is the Einstein A
coefficient, ν is the frequency of the J=u-ℓ transition, and
TMBdv is the integrated intensity
of the line. The column density of this one state can then be related to the total column
density of that molecule through the partition function. It is the total column density for
the molecule (not the observed state) that is presented in Table 3.
3.1. Source Properties derived from SiO observations
For each source we determined the column density, or upper limit to the column density,
in SiO (Table 3) using the methods described above. These column densities can be compared
to the column densities of other molecules (i.e. CS) for the same regions in order to determine
the fractional abundance of SiO, if the abundance of the other molecule is known. We were
able to obtain the CS or C34S column densities for fifteen of our sources from the literature.
Column densities for fourteen sources were taken from Plume et al. (1997, hereafter P97),
with the column density for one additional source taken from Wang et al. (1993). For
those sources with C34S column densities instead of CS column densities, we assumed an
abundance ratio of [CS]/[C34S] = 22 (Wilson & Rood, 1994) to determine a CS column
density. The abundance of CS, relative to H2, was calculated by S03 for 13 of these sources,
and we assume a CS abundance of 1.2 × 10−9 for the other two source for which the CS
column density is known, since this was the average CS abundance as calculated by S03.
We then compare the column density and abundance of CS to our observed SiO column
density, or column density upper limit, to determine the abundance of SiO relative to H2 in
our sources (Table 3).
Despite our source selection criteria requiring previous evidence of outflows, we detected
SiO towards only 14 of our 23 sources. This raises a number of questions, such as: is the
observed SiO in fact tracing outflow if we do not detect it in all sources? Why do we not
detect SiO in all sources? Is the signal being beam diluted at large distances? Has the Si
evolved into other species? Below we first address whether the detected SiO can be used
as an outflow tracer, and then discuss reasons for our non-detections of SiO in nine of our
sources.
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Si is liberated in shocks, and if these shocks are not due to the outflow, they must be due
to the photo dissociation region (PDR) surrounding the UCHII region. Evidence for diffuse
(not collimated) SiO can be seen in W75N (Shepherd, Kurtz & Testi, 2004) suggesting that
the SiO may be due to the PDR and not an outflow. Models and observations of SiO in
the PDRs around high mass star forming regions suggest moderate SiO enhancement, and
that the SiO abundance is independent of the ambient radiation field (i.e. Schilke et al.
2001). Schilke et al. (2001) find SiO column densities of ∼ 1012 cm−2 in their observed
PDRs. This is, admittedly, below our detection threshold; however, we detect average SiO
column densities of ∼ 1014 cm−2. This suggests a possibly higher SiO abundance than found
in PDRs.
The enhanced SiO column density alone is not enough to discount the origin of the SiO
in our sources as the PDR and so we can consider how our SiO abundance varies with the
ambient radiation field. For the gas near an HII region, we can approximate the strength
of the ambient radiation field using the Far Infrared (FIR) luminosity of the region. For
twelve of our sources with SiO detections and abundance calculations, we obtained the FIR
luminosity from either Wood & Churchwell (1989), Kurtz et al. (1994) or Evans et al. (1981).
We then scaled their values for the different source distances used in this study (see Table 3).
Comparing the SiO abundance to the source luminosity (see Figure 5), we find that the SiO
abundance increases with source luminosity. There is a 10% chance that this relationship
could arise from uncorrelated data. Thus, it is possible that our result is contrary to the
findings of Schilke et al (2001), and we suggest that the SiO we observe does come primarily
from outflow shocks.
We can also compare our detection rate of SiO to that of the SiO survey towards maser
sources of Harju et al (1998) who observed SiO (J=2-1) and SiO (J=3-2). For comparison
to our results, we only consider the sources in Harju et al. which are listed as UCHII
regions. Our detection rate is 61%, compared to their rate of 29%. While our detection
threshold is slightly lower (our observations have rms noise levels generally below 0.05 K
at 347 GHz, while their rms noise levels are generally below 0.08 K), we suggest that the
different detection rates are due to differences in the source selection criteria. Although both
samples contain UCHII regions, our sample contains sources with previous observations of
outflows, while Harju et al. have selected sources based on previous observations of masers.
For the 12 sources which overlap between the two studies, we detect SiO towards 9 sources,
while they detect SiO towards 10. They detected SiO (J=2-1) in G31.41 while we did not
detect it in SiO (J=8-7).
Based on this comparison to SiO observations of UCHII regions not selected by outflows,
which have a lower SiO detection rate, and that our SiO abundance increases with source
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luminosity, we suggest that the SiO, in the 14 sources in which it is detected, is being
generated in the outflow. Previous, high resolution observations of SiO also suggest that
SiO can be enhanced in the outflows from high mass stars (i.e. Beuther et al. 2004, Beuther,
Schilke & Gueth 2004) just as it is in the outflows from low mass stars.
The nine non-detections in our sample could be caused by beam dilution if these sources
are on average further away. However, if we compare the distances for sources with and
without SiO detections, we find average distances of 6.3 ± 4.4 and 4.4 ± 2.5 kpc, respectively.
Thus, the non-detections cannot be attributed to larger average source distances and so beam
dilution can play only a minimal role in the non-detections of SiO.
If these SiO non-detections are not due to distance effects, there must be some local
phenomenon which can explain why SiO is not being detected in regions known to contain
protostellar outflows. It is possible that the Si is evolving into different species and the SiO
abundance is dropping back down to dark cloud values. Pineau des Forêts et al. (1997)
suggest that a few×104 yr after the Si is liberated from dust grains and forms SiO, it can
either freeze out back onto dust grains or oxidize and form SiO2. Thus, the lack of SiO may
be due to silicon moving into other species if it was liberated more than 104 yr ago. This
interpretation implies that we did not detect SiO in some of our sources because the outflow
generating mechanism shut off more than 104 years ago, and the outflow observed in HCO+
(or in other molecules by other authors) is a remnant of previous accretion.
The kinematic ages of nine of our sources are listed in the Wu et al (2004) catalog of high
velocity outflows. Of these nine sources, five were also included in the P97 and S03 studies.
This results in five sources for which we have both the kinematic age of the outflow, and
the abundance of SiO. The relationship between outflow age and SiO abundance is shown
in Figure 6 along with the model predictions of Pineau des Forêts et al. (1997). With only
five points, it is difficult to draw conclusions about the relationship between outflow age
and SiO abundance, especially given the uncertain beam filling factor. At higher resolution,
these points would likely move upwards to higher abundances. A larger beam filling factor
would move the points for the two young outflows (G5.89 and Cep A) towards the model
predictions. As for the other three outflows, this would move them further from the model
predictions. We suggest that the outflow generating mechanism is continuing to shock these
regions, replenishing the SiO. The oldest of these sources (G192.58) shows an infall signature
in HCO+, which also suggests that the outflow is still being powered.
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3.2. Source properties derived from HCO+ observations
Given the large average distance to our sources (5.7 kpc), our 15′′ beam subtends an
average linear distance of 0.4 pc. Thus, it is quite likely that the HCO+ emitting region does
not entirely fill our single JCMT beam. In general we can determine the beam filling factor,
f , for each source using
TL = fsourceTs(1− e−τ ) (2)
where TL is the line brightness temperature measured at the telescope (corrected for telescope
efficiencies) and Ts is an approximation to the ambient temperature (44 K) which is valid
at densities greater than ∼ 103 cm−3 (Rholfs & Wilson, 1994). In a number of our sources,
the HCO+ is asymmetric, and thus cannot be consistently used to determine a beam filling
factor. Instead, we can use the optically thin H13CO+ line, and we can simplify the above
equation to f = TL/(Tsτ). These values are shown in Table 3.
HCO+ becomes optically thick quite quickly due to its relatively high abundance with
respect to H2 ([HCO
+]/[H2] ∼ 10−8), and as such, can be used to roughly trace outflow and
to trace infall (i.e. Myers et al. 1996) if the line profile shows a double peak. We determined
the optical depth of HCO+ towards each source using Equation 1 of Choi et al. (1993), and
found that in all but one case (G139.9), it is optically thick. In all cases, the optical depth
of HCO+ is less than 77 (the abundance ratio between HCO+ and H13CO+, Wilson & Rood,
1994), resulting in optically thin H13CO+ towards all sources. Because HCO+ is optically
thick in its line center, the line wings can be used to detect outflows, and so, if there is
an outflow, it should be detectable in HCO+ even if it goes undetected in SiO. Gaussian
profiles were fit to our HCO+ spectra (either single or double Gaussians, depending on the
observed line shape), and the fits were subsequently subtracted from the spectra to leave
only the residual outflowing gas. When using two Gaussians to fit the self absorbed spectra,
we employed a method similar to the single Gaussian fitting of Purcell et al. (2006) because
we used the sides of the detected lines to fit our profiles (see their Figure 3). Comparing the
two Gaussian fits to single Gaussian fits showed no significant differences in distinguishing
line wing intensity. Because of the possibility of contamination from SO2 emission in the
blue shifted outflow wing (at -17 km s−1), the peak brightness of the residual emission was
determined using only the red shifted wing emission. In all cases (except for G10.47 for
which we could not find a linear baseline), we found a minimum of a 5σ peak brightness
temperature in the residual line wing emission, with 19 of our sources having a minimum of
a 10σ peak. This result suggests that we can detect outflow motions in all sources using our
detections of HCO+, despite not detecting SiO towards every source.
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Our observations show that for ten of our sources, the spectral line profile of the HCO+
emission has a double peak. This profile could either be due to self absorption of the optically
thick HCO+ line or from multiple velocity components within our 15′′ beam. To break this
degeneracy, we observed the optically thin H13CO+. If the H13CO+ line has a single peak at
the same velocity as the HCO+ absorption feature, then it is likely that the HCO+ line is self
absorbed. If, however, the H13CO+ also has two peaks, and they are at approximately the
same velocities as the two HCO+ peaks, it would suggest that there are multiple components
within the beam. Of our ten sources with double HCO+ peaks, only one shows a double
peak profile in H13CO+ (G20.08). This results in nine sources with double peaked optically
thick HCO+.
In addition to the nine optically thick sources, similar line asymmetries appear in a
number of other sources. However, in these sources, there is no clear emission gap producing
a double peak profile, only an emission shoulder (i.e. De Vries & Myers 2005). If we take
G75.78 as an example, the HCO+ line peak is red shifted from the rest velocity of the source,
with a blue shifted emission shoulder.
There are a number of different kinds of source dynamics that can lead to the double
peaked line profiles seen in our spectra, such as infall, outflow and even rotation. However,
infall is the only one of these processes which would produce line asymmetries which are
consistently blue (i.e. the blue peak is higher than the red peak or shoulder). If these
profiles were due to outflow or rotation, there would be no statistical reason to have more
sources with higher blue peaks than red peaks. Many previous studies have investigated the
statistical significance of using this type of optically thick blue line asymmetry to trace infall
as opposed to other dynamical motions (i.e. Mardones et al. 1997 and Gregersen et al. 1997
for low mass star forming regions, and Fuller et al. 2005 for high mass star forming regions).
Of the 10 sources in our survey which have double peaked HCO+ profiles, we suggest
eight may be indicative of infall. The other two sources are G20.08 and G45.47. G20.08
has already been shown to have multiple components in the beam from the double peaked
H13CO+ profile, and G45.47 has a brighter red peak than blue. There are two additional
sources (G19.61, and G240.3) in which HCO+ has a strong red shifted shoulder, which we
suggest may also be tracing infall. This analysis gives a total of ten infall candidates in our
sample of 23 sources.
A recent survey of HCO+ (J=1-0) towards sources with methanol masers shows an even
distribution of sources with blue and red line asymmetries, and a higher percentage of self
absorbed lines than in our study (Purcell et al. 2006). Of the six sources which overlap
between our survey and that of Purcell et al, all six are self absorbed in HCO+ (J=1-0).
Five of them have blue line asymmetries consistent with infall, while only one (G31.41) has
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its red peak brighter than its blue peak. We only find self-absorption in HCO+ (J=4-3) for
three of these six sources. In two of the sources for which we do not see a clear self-absorption
feature, we do see evidence for a red shifted shoulder which may be showing unresolved infall.
The sources in Purcell et al. (2006) have an even distribution of red and blue line
asymmetries, while we have a clear bias towards detecting blue line asymmetries. This
comparison could suggest that the higher energy J=4-3 transition of HCO+ is a better
tracer of infalling gas because it does not self absorb as readily as the J=1-0 transition.
For each of the 8 sources with blue, double peaked HCO+ profiles, we can determine an
infall velocity (vin) using the two layer radiative transfer model of Myers et al. (1996). Using
their equation 9, we find infall velocities for all eight double peaked infall sources (Table 4).
The mass infall rate can then be determined using:
ρV vin
πnH2µmHr
gmvin (3)
where µ is the mean molecular weight (µ = 2.35), the geometric mean radius (rgm) is the
unresolved circular radius of the HCO+ emitting region derived from the beam radius and
the beam filling factor (rgm =
frbeam), and nH2 is the ambient source density. For seven
of the sources in Table 4, the ambient density was determined by either P97, Hofner et
al. (2000), or Wang et al. (1993). We could not find the ambient density for the eighth
source (G192.58). From this analysis, we determine mass infall rates ranging from 1×10−2
to 2×10−5 M⊙ yr−1. These values are slightly higher than those generally observed for low
mass star forming regions, but are consistent with the accretion rates derived for high mass
star forming regions by McKee & Tan (2003). Since outflow rates are orders of magnitude
higher in high mass star forming regions (i.e. Beuther et al. 2002) it is not unreasonable to
suggest that infall rates are also much higher in these regions.
The mass outflow rate for only one of these sources (Cep A) can be determined from
the Wu et al. (2004) survey of high velocity outflows by dividing the mass in the outflow
by the kinematic age of the outflow. We find the ratio of the mass outflow rate to the mass
infall rate to be Ṁout/Ṁin ≈ 16. This value is only slightly higher than values seen in other
high mass star forming regions (i.e. Behrend & Maeder 2001). Also, models suggest that a
mass equivalent to 20-30% of the mass accreted onto a protostar is ejected as a wind (i.e.
Pelletier & Pudritz 1992, Shu et al. 1994), and that this wind entrains 5-20 times its mass
in the outflow (Matzner & McKee 1999).
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3.3. Source properties derived from mapping G45.07
At a distance of 9.7 kpc, G45.07 is one of our furthest sources. This source was known
from previous observations to have multiple continuum sources (De Buizer et al. 2003,
2005). The three continuum sources were observed in the mid-Infrared (MIR) and all three
fall within 6′′ of our map center. A fairly young outflow has also been mapped at high
resolution (< 3′′ synthesized beam) in CO and CS towards this region (Hunter et al. 1997).
They observed a bipolar outflow with a position angle of -30◦ (east of north), as well as a
red shifted absorption feature in their CS observations which they take to be indicative of
infall.
Due to the large distance to this source, we should be able to detect all of the emission
associated with this source in a fairly small map. In the left panel of Figure 7 we present a
map of the SiO (contours) and H13CO+ (halftone) emission in this region. The right panel
of Figure 7 shows the HCO+ emission. In both figures, the 5σ (2.1 K km s−1) H13CO+
emission contour is plotted as a dashed line to help guide the eye. The first contour for SiO
in the left panel and only contour of HCO+ in the right panel are also 5σ (2.2 and 3.5 K
km s−1 respectively). The differences in the 5σ contour levels for each tracer come from the
different single channel rms levels between the two tunings, and the width of each line as
given in Table 2. Also plotted in both figures are the three MIR continuum sources observed
by De Buizer et al (2003, 2005).
If we did not have the added information provided by this map, there would be two
main conclusions we could draw from our single pointing observations towards this source.
The first is that the enhanced blue emission in all three tracers suggests that our pointing
is observing more of the blue shifted outflow lobe than the red. Second, since we only have
one peak in the spectrum of each tracer, there is only one source and we cannot classify it
as infalling.
The left panel of Figure 7 shows contours of SiO emission superimposed on the H13CO+
halftone. With beam spacings of 5′′, these maps are oversampled; however we note that much
of the structure in the SiO emission is on scales comparable to the size of the JCMT beam.
For instance, the structure at ∆α = −10′′, ∆δ = 0′′ is offset from the map center by more
than the radius of our beam and could be independent from the emission at the map center.
There is also SiO emission at ∆α = 5′′, ∆δ = 15′′, which is more than a full beam away
from the map center, and suggests that the SiO emission is more extended that the primary
beam of our observations. In fact, it appears as though there is a second SiO emission peak
towards the upper left of the left panel of Figure 7. Interestingly, there does not appear to
be as much H13CO+ emission at this northern position. This comparison shows that the
SiO and H13CO+ lines are tracing different gas populations in this region. The excess SiO
– 13 –
emission is offset from the map center in the same direction as the CO emission shown in
Hunter et al. (1997) at much higher resolution.
The right panel of Figure 7 shows the HCO+ emission for this region. It appears that
the HCO+ emission extends much further than the SiO emission, suggesting it is tracing the
larger scale envelope material. The line through the middle of this plot indicates the cut
taken for the position-velocity (PV) diagram along the outflow axis as described by Hunter
et al. (1997)
The two panels of Figure 8 show the PV diagrams for SiO and HCO+ in our maps both
perpendicular and parallel to the outflow axis defined by Hunter et al. (1997). Our single
pointing HCO+ spectrum (Figure 4) suggests we are observing more blue shifted outflow
emission than red shifted emission; however, from our PV diagrams, we see that there is
excess blue emission at all positions in our map. This excess blue emission cannot be due
to outflow alone; instead, it could be due to an inherent velocity shift between the three
continuum sources in our beam. We can, in fact, fit three Gaussian components to most of
our HCO+ spectra. These Gaussians peak at velocities of 60, 52 and 44 km s−1, with the
peak temperature for each component decreasing with velocity. The third component (at 44
km s−1) could not be fit at all positions because it was intrinsically weaker than the other
two peaks, and was lost in the noise towards the edges of the map. It appears as though
this third component might be contamination from SO2, which should occur at an apparent
velocity of 41 km s−1 (or -17 km s−1 in Figure 4).
Perpendicular to the outflow axis, the mean velocity of the HCO+ line peak appears
to shift from ∼ 58 km s−1 at an offset of +15′′ from the source center to ∼ 61 km s−1 at
an offset of −15′′ from the source (Figure 8). Given our velocity resolution (1.08 km s−1)
and spatial resolution (15′′), is unclear whether this velocity shift is real. If it is, it could
indicate large scale (∼ 1.4 pc) rotation within the core, on a much larger scale than would
be expected for a rotating accretion disk.
3.4. Correlations between Datasets
Previously, we discussed the reasons why we do not detect SiO towards a number of
our sources, and have calculated the mass infall rates for the sources with double peaks
in their HCO+ emission, but we have not yet discussed the correlations between the two
species. In Figure 9 we plot the logarithm of the abundance of H13CO+ against the same
quantity for SiO (open circles), as well as the column densities of both species (filled circles).
The abundance of H13CO+ was calculated in the same manner as the abundance of SiO
– 14 –
described above (using the CS column density from P97 and the CS abundance from S03).
The probability of obtaining these correlations if the data are, in fact, uncorrelated is 6×10−3
for the abundances, and 2×10−4 for the column densities.
As stated earlier, SiO is a well known shock tracer, and as such, an increased abundance
of SiO would suggest more shocked material within our 15′′ beam. The (generally) infall
tracing HCO+ has been shown by some authors to be destroyed in strong shocks (i.e. Bergin
et al. 1998, Jørgensen et al. 2004). However, Wolfire & Königl (1993) suggest that HCO+
can be enhanced in regions with high energy shocks, where electron abundances are much
greater. This enhancement in the electron abundance increases the formation rate of ions,
and we suggest that this is responsible for the H13CO+ abundance enhancement in our
sources. This correlation between the abundances of H13CO+ and SiO suggests that HCO+
and H13CO+ are not only tracing infalling gas, but also the outflowing gas as well. This
conclusion is supported by the strong, and broad, line wing emission detected in HCO+ (See
Section 3.2).
HCO+ over abundances have been seen in high mass star forming regions not included
in this study like NGC 2071 (Girart et al. 1999) and Orion IRc2 (Vogel et al. 1984). In
these two papers, the over abundances of H13CO+ are with respect to ambient cloud tracers
such as CO and H2, rather than the high density or shock tracers like the CS and SiO with
which we are comparing our H13CO+ abundances. However, Viti & Williams (1999a,b) show
that HCO+ is indeed over abundant with respect to CS in the gas surrounding HH objects,
and Jiménez-Serra et al. (2006) also show that the abundance of H13CO+ can be enhanced
with respect to SiO by up to a factor of ten in the same regions ahead of HH objects.
4. Discussion and Conclusions
Without maps of each region, it is impossible to tell how much of the HCO+ emission in
the line center and in the line wings is due directly to infall and outflow motions; however,
based on the arguments we have presented above, we suggest that ten sources show infall
motions, and all 23 source show outflow motions based on the HCO+ line profiles. We have
found evidence for recent outflow activity (SiO emission) in 14 out of our 23 sources. Seven
of these outflow and infall sources overlap. M17S, G192.6 and G240 appear to show only
infall signatures and no SiO outflow signatures. They do, however, appear to have HCO+
outflow signatures of a minimum of 8σ.
Detection of line wing emission in HCO+ and the relationship between H13CO+ and
SiO abundances described in the previous section suggest that while SiO is tracing outflow
– 15 –
in most sources and HCO+ is tracing infall in some sources, HCO+ is also observable in the
outflowing gas for all regions.
We find that the non-detection of SiO in nine of our sources is not due to beam dilution
or larger average distances to the source, but possibly to older outflows for which the Si has
likely either frozen back onto dust grains or evolved into SiO2. For these sources, it appears
that the accretion may have ceased, and the observed outflowing gas is a remnant of previous
accretion.
We have found seven sources with SiO outflow signatures but no infall signatures in
HCO+. This result could be due to a number of factors such as beam dilution of the
infalling gas which masks the spectral line profile we would expect for large scale infall.
It is possible that, as the outflow ages and widens, it may impinge on the region in
which we could detect infalling gas. For outflow cones oriented along the line of sight, the
younger, narrower outflows would have infalling gas with large line of sight velocities and be
likely to produce an observable infall signature. However, for the older outflows which have
widened, the largest infall velocities will be in the plane of the sky, and unobservable at the
resolution of the JCMT. However, this effect would not be as pronounced for outflows in the
plane of the sky (such as G5.89, for which we do not see an infall signature). It is difficult to
assess the importance of this effect without detailed information on the outflow orientation
in each source.
Thus, we suggest that some of these sources may have finished accreting, and what we
observe are remnant outflows from a previous phase of accretion. This scenario was suggested
by Klaassen et al. (2006) to explain the large scale outflow in G5.89, and this source is one
of these seven sources with an SiO outflow and no apparent infall signature.
The seven sources which show recent outflow activity (those with SiO emission) and
which appear to be undergoing infall are suggestive of ongoing accretion beyond the onset of
the HII region. If accretion is ongoing in the presence of an HII region, then it seems likely
that this accretion flow may be ionized. This ionized accretion scenario could be similar to
low angular momentum accretion with high ionization as suggested by Keto (2006). Thus,
from this survey, we suggest that both models of ionized accretion and halted accretion may
be important in describing the evolution of a massive protostar (or protostars) beyond the
formation of an HII region.
We would like to acknowledge the support of the National Science and Engineering
Research Council of Canada (NSERC). We thanks the referee for helpful comments which
improved the paper. P.D.K would also like to thank E. Keto and D. Johnstone for helpful
– 16 –
discussions during the preparation of this manuscript.
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– 20 –
Table 1. Observed Sample of Massive Star Forming Regions.
Name Position (J2000) RMS noise limit (K) VLSR Distance P92 Name
RA DEC 347 GHz 356 GHz (km s−1) (kpc) ref
G5.89 18 00 30.3 -24 03 58 0.060 0.111 9 2 1 W28A2 (1)
G5.97 18 03 40.4 -24 22 44 0.044 0.069 10 2.7 5 · · ·
G8.67 18 06 19.0 -21 37 32 0.068 0.074 36 8.5 2 8.67-0.36
G10.47 18 08 38.4 -19 51 52 · · · b 0.100 67 12 2 W31 (1)
G12.21 18 12 39.7 -18 24 21 0.042 0.076 24 16.3 2 12.21-0.1
M17S 18 20 24.8 -16 11 35 0.044 0.077 20 2.3 5 M17 (2)
G19.61 18 27 38.1 -11 56 40 0.048 0.065 43 4.5 1 19.61-0.23
G20.08 18 28 10.4 -11 28 49 0.044 0.068 42 4.1 1 20.08-0.13
G29.96 18 46 03.9 -02 39 22 0.044 0.079 98 9 1 W43S
G31.41 18 47 33.0 -01 12 36 0.044 0.061 97 8.5 1 31.41+0.31
G34.26 18 53 18.5 01 14 58 0.047 0.131 58 3.7 1 W44
G45.07c 19 13 22.1 10 50 53 0.044 0.073 59 9.7 1 45.07+0.13
G45.47 19 14 25.6 11 09 26 0.037 0.063 58 8.3 6 · · ·
G61.48 19 46 49.2 25 12 48 0.037 0.048 12 2 1 S88 B
K3-50A 20 01 45.6 33 32 42 0.035 0.077 -24 8.6 3 K3-50
G75.78 20 21 44.1 37 26 40 0.038 0.068 0 5.6 1 ON 2N
Cep A 22 56 17.9 62 01 49 0.031 0.076 -10 0.7 1 CEP A
W3(OH) 02 27 03.8 61 52 25 0.026 0.079 -48 2.4 2 W3 (OH)
G138.3 03 01 29.2 60 29 12 0.066 0.073 -38 3.8 1 S201
G139.9 03 07 23.9 58 30 53 0.058 0.071 -39 4.2 1 · · ·
G192.58 06 12 53.6 17 59 27 0.037 0.053 9 2.5 3 S255/7
G192.6 06 12 53.6 18 00 26 0.074 0.098 9 2.5 3 S255/7
G240.3 07 44 51.9 -24 07 40 0.035 0.048 68 6.4 4 · · ·
aName given to source in Plume et al. (1992). When included in P97 and S03, these were the
source names used.
bToo much chemistry in spectrum to determine a reliable rms noise level.
– 21 –
cThe rms noise limits for the two maps of this source are 0.1 K and 0.14 K for the 347 GHz
and 356 GHz pointings respectively.
Note. — The source velocity here is the reference velocity of each source, which has been
removed from each spectrum. RMS noise levels are given in TMB. References:(1) Hanson et al.
2002, (2) S03, (3) Kurtz et al. 1994, (4) Kumar et al. 2003., (5) Churchwell et al. 1990, and (6)
Hofner et al. 2000
– 22 –
Table 2. Integrated Intensities for SiO, HCO+ and H13CO+.
Name SiO Properties HCO+ Properties H13CO+ Properties
TMBdv dv TMB
TMBdv dv TMB
TMBdv dv
G5.89 2.3 75.3± 0.7 120 39.2 690.0±0.8 105 8.9 65.2± 0.4 35
G5.97 <0.1 <0.3± 0.1 8 12.4 50.5 ±0.2 15 0.6 1.6 ± 0.1 8
G8.67 0.2 1.6 ± 0.3 13 13.8 80.0 ±0.3 38 5.6 20.9± 0.2 10
G10.47 1.0 9.9 ± · · · 22 7.7 79.3 ±0.4 35 1.2 9.2 ± · · · 15
G12.21 0.2 2.5 ± 0.2 25 10.0 93.8 ±0.3 30 0.8 6.5 ± 0.2 15
M17S <0.1 <0.7± 0.2 15 16.2 79.0 ±0.2 15 2.9 9.4 ± 0.2 11
G19.61 0.8 15.9± 0.3 45 10.8 162.2±0.4 55 1.8 16.1± 0.2 20
G20.08 0.4 6.9 ± 0.3 40 6.5 75.0 ±0.3 40 1.1 9.7 ± 0.2 25
G29.96 0.5 7.9 ± 0.3 40 23.6 160.4±0.3 30 3.8 13.2± 0.1 10
G31.41 <0.1 <0.7± 0.2 15 1.6 6.8 ±0.2 13 <0.1 <0.3± 0.1 3
G34.26 1.3 26.0± 0.4 55 27.8 228.2±0.6 35 4.3 33.2± 0.2 25
G45.07 0.7 10.3± 0.3 37 15.0 169.4±0.4 45 1.5 11.2± 0.3 33
G45.47 <0.1 <0.3± 0.1 12 10.1 59.4 ±0.2 20 1.3 6.1 ± 0.1 12
G61.48 <0.1 <0.3± 0.1 8 10.8 58.5 ±0.1 18 1.5 4.7 ± 0.1 8
K3-50A 0.3 2.1 ± 0.1 16 18.6 154.8±0.3 25 1.8 10.6± 0.2 18
G75.78 0.3 2.2 ± 0.1 13 12.5 136.8±0.3 45 2.5 10.3± 0.2 17
Cep A 0.5 7.3 ± 0.2 45 23.2 234.7±0.4 60 4.5 23.4± 0.2 28
W3(OH) 1.4 14.5± 0.2 50 18.5 148.2±0.3 30 2.4 11.6± 0.1 15
G138.3 <0.2 <0.3± 0.1 2 5.0 10.6 ±0.2 8 0.3 0.5 ± 0.1 4
G139.9 <0.2 <0.3± 0.1 11 8.6 18.8 ±0.1 6 <0.2 <0.3± 0.1 2
G192.58 0.3 1.3 ± 0.1 8 14.5 84.2 ±0.2 25 1.1 5.5 ± 0.1 15
G192.6 <0.2 <0.6± 0.2 30 12.0 54.1 ±0.3 15 0.8 2.2 ± 0.2 8
G240.3 <0.1 <0.4± 0.2 10 7.3 59.1 ±0.2 30 0.4 2.5 ± 0.1 14
Note. — For the sources in which we did not detect SiO or H13CO+, 3σ upper limits
on the brightness temperature and integrated intensities are given.
– 23 –
Table 3. Observed and Derived Source Parameters.
Name Detection f a Column density (×1012) [SiO]/[H2]b LFIR toutflowc
SO2 Infall H
13CO+ SiO Log(L⊙) (10
4 yr)
G5.89 Y N 0.79 39.7 421.0 -8.90 5.25d 0.2
G5.97 N N 0.28 1.0 <1.7 · · · 5.23d · · ·
G8.67 Y Y 0.24 12.7 9.0 · · · 5.70d · · ·
G10.47 Y Y 0.16 5.6 55.4 -9.53 6.26d · · ·
G12.21 Y N 0.22 4.0 14.0 -10.15 6.17d · · ·
M17S N Y 0.33 5.7 <3.9 <-10.73 5.72d · · ·
G19.61 Y Y 0.22 9.8 89.0 -9.72 5.42d · · ·
G20.08 Y N 0.13 5.9 38.6 -10.48 4.86d · · ·
G29.96 Y N 0.49 8.0 44.2 -9.12 6.30d · · ·
G31.41 N N · · · <0.2 <3.9 <-11.89 5.45d · · ·
G34.26 Y Y 0.58 20.2 145.0 -10.43 5.77d · · ·
G45.07 Y N 0.32 6.8 57.6 -9.54 6.15d 4
G45.47 Y N 0.21 3.7 <1.7 · · · 6.04d · · ·
G61.48 N N 0.23 2.9 <1.7 <-10.80 5.01d 7
K3-50A Y N 0.40 6.5 11.7 -11.60 6.35e · · ·
G75.78 Y N 0.25 6.3 12.3 · · · 5.65d 3.7
Cep A Y Y 0.47 14.2 40.8 -9.78 4.40f 0.2
W3(OH) Y Y 0.39 7.1 81.1 -11.16 5.12e · · ·
G138.3 N N 0.11 0.3 <1.7 · · · 4.57e 17
G139.9 N N · · · <0.2 <1.7 · · · 4.82e 6
G192.58 Y Y 0.32 3.4 7.3 -10.94 4.79e 50
G192.6 Y Y 0.26 1.3 <3.4 · · · · · · · · ·
G240.3 Y Y 0.16 1.5 <2.2 · · · · · · 2.3
aHCO+ beam filling factor.
bSiO abundance relative to H2 for sources with CS observations (P97) and abundance
calculations (S03).
cKinematic age of outflow from Wu et al. (2004).
– 24 –
dFar Infrared Luminosities modified from Wood & Churchwell (1989).
eFar Infrared Luminosities modified from Kurtz et al. (1994).
fFar Infrared Luminosities taken from Evans et al. (1981).
– 25 –
Table 4. Infall Velocities and Mass Infall Rates.
Name Vin
a nH2
b Ṁin
(km s−1) (105 cm−1) (10−4 M⊙ yr
G8.67 0.4 ± 0.1 1.8 4 ± 2
G10.47 1.8 ± 0.3 7.2 100± 80
M17S 1.4 ± 0.5 5.0 4 ± 2
Cep A 0.23±0.07 10.0 0.17± 0.07
W3(OH) 0.06±0.02 60.0 3 ± 1
G192.58 0.8 ± 0.3 · · · · · ·
G192.6 0.9 ± 0.4 4.0 2 ± 1
G34.26 1.5 ± 0.3 3.6 14 ± 4
aInfall velocities for sources with double peaked
HCO+ profiles.
bAmbient densities taken from P97 except for:
G34.26 (Hofner et al. 2000), M17S (Wang et al. 1993).
cMass infall rates.
– 26 –
Fig. 1.— Nine sources in which SiO was not detected (to 4 σ limits). The name of each
source is given in the top left hand corner of each panel. For each source, the source rest
velocity is plotted as ∆v = 0 km s−1.
– 27 –
Fig. 2.— HCO+ and H13CO+ sources with no SiO detections (the same sources as in Figure
1). Solid lines show HCO+ emission, while dashed lines show H13CO+ emission scaled up by
a factor of four. The temperature scale on the left hand side of the panels is the scale used
for the HCO+ spectra, and the temperature scale on the right hand side is that used for the
H13CO+ spectra. For sources with SO2 detections, we have re-plotted the HCO
+ spectra on
a larger intensity scale ( from -1 K to 3 K, in gray) to show the low lying SO2 emission. The
SO2 emission, when present, is centered at -17 km s
−1, and is indicated by an arrow. G31.41
is slightly offset to emphasize that the temperature scale has been magnified to show the
emission feature.
– 28 –
Fig. 3.— SiO spectra for 14 source in which SiO was detected (above 4 σ limits). For this
figure only, the range of temperatures plotted increases by row. The plotted temperature
ranges for each row increase by 0.5 K per row. The temperature scale in the top row extends
from 0.2 to 1 K (in the TMB scale), while the temperature scale in the bottom row extends
from 0.2 to 2.5 K. Note that the velocity scale for G5.89 is much wider (a 78 km s−1 window
as opposed to a 38 km s−1 window for the other sources) in order to show the full width of
the emission line, but that the temperature scale is the same as that for W3(OH).
– 29 –
Fig. 4.— HCO+ and H13CO+ emission from sources with SiO detections (the same sources
as in Figure 3). The solid, dashed and gray lines are the same as shown in Figure 2, as is
the placement of the temperature scales. As in Figure 3, G5.89 has been plotted separately
from the rest of the sources to stress that the velocity scale is larger for this source.
– 30 –
Fig. 5.— SiO abundance plotted as a function of source luminosity. The line of best fit shown
is given by (Log([SiO]/[H2])=0.30±0.06Log(L/L⊙)-11.7±0.3), and was calculated only for
sources detected in SiO. This shows that SiO abundance increases with source luminosity.
This is contrary to what is expected for SiO produced in PDRs, suggesting the observed SiO
is produced in outflows.
– 31 –
Fig. 6.— SiO abundance plotted as a function of outflow age. The outflow ages were taken
from Wu et al. (2004). The solid line shows the model predictions of Pineau des Foréts et al.
(1997) for the SiO abundance as a function of age (assuming Si returns to the dust grains),
and the dashed line represents the canonical dark cloud abundance of SiO. The source shown
with a downwards arrow represents G61.48, a source in which we did not detect SiO, and the
value given is an upper limit to the SiO abundance. The error bars represent 30% calibration
uncertainty between our observations and those of P97 and S03
– 32 –
Fig. 7.— Left: SiO and H13CO+ towards G45.07. The halftone scale represents the in-
tegrated H13CO+ emission, while the contours show the integrated intensity of SiO. The
velocity range used to determine the integrated intensities is the same as the single pointing
velocity range (dv), however the rms noise limits are those listed at the end of Section 2.
The first SiO contour is 5σ (3.2 K km s−1), incrementing in steps of 2σ, with the same
scale continuing into the white contours near the center. Right: HCO+ integrated intensity
towards G45.07. The solid contour represents the 5σ emission for the HCO+ emission (3.5
K km s−1), and the solid line represents the cut used for the PV diagram along the outflow
axis (PA = -30◦) presented in Figure 8. For both panels, the dashed line represents the 5σ
H13CO+ emission contour (3.0 K km s−1). The three triangles represent the Mid IR sources
detected by De Buizer et al. (2005), with the two points in the center corresponding to 230
GHz continuum sources detected at the Submillimeter Array (Klaassen et al. in prep).
– 33 –
Fig. 8.— Position-Velocity (PV) diagrams for SiO (halftone scale) and HCO+ emission
(contours) in G45.07 both perpendicular to (top panel) and along the outflow axis (bottom
panel) as defined in Hunter et al. (1997). These PV diagrams were taken at positions angles
of 60◦ and -30◦ east of north respectively, with the cut for the bottom panel of this figure
shown in the right panel of Figure 7. The first contour for HCO+ is 5σ, or 0.7 K since the
rms noise limit for this map is 0.14 K, and the contours increase in increments of 5σ. The
three dashed lines in the bottom panel show the peak velocities of the Gaussian fits to the
HCO+ spectra.
– 34 –
Fig. 9.— The abundance of H13CO+ with respect to the abundance of SiO (open
circles, dashed line of best fit) appears to increase faster than the respective column
densities of these two species (filled circles, solid line of best fit). The equations of
the two lines of best fit are: Log[X(H13CO+)]=(0.84±0.09)Log[X(SiO)]-(2.4±0.9), and
Log[N(H13CO+)]=(0.5±0.1)Log[N(SiO)]+(5±1). This relationship could be due to HCO+
being enhanced (similarly to SiO), as discussed in the text.
	Introduction
	Observations
	Results
	Source Properties derived from SiO observations
	Source properties derived from HCO+ observations
	Source properties derived from mapping G45.07
	Correlations between Datasets
	Discussion and Conclusions

  We present single pointing observations of SiO, HCO$^+$ and H$^{13}$CO$^+$
from the James Clerk Maxwell Telescope towards 23 massive star forming regions
previously known to contain molecular outflows and ultracompact HII regions. We
detected SiO towards 14 sources and suggest that the non-detections in the
other nine sources could be due to those outflows being older and without
ongoing shocks to replenish the SiO. We serendipitously detected SO$_2$ towards
17 sources in the same tuning as HCO$^+$. We detected HCO$^+$ towards all
sources, and suggest that it is tracing infall in nine cases. For seven infall
candidates, we estimate mass infall rates between 1$\times10^{-2}$ and
2$\times10^{-5}$ M$_{\odot}$ yr$^{-1}$. Seven sources show both SiO detections
(young outflows) and HCO$^+$ infall signatures. We also find that the abundance
of H$^{13}$CO$^+$ tends to increase along with the abundance of SiO in sources
for which we could determine abundances. We discuss these results with respect
to current theories of massive star formation via accretion. From this survey,
we suggest that perhaps both models of ionized accretion and halted accretion
may be important in describing the evolution of a massive protostar (or
protostars) beyond the formation of an HII region.

Introduction
The dynamics in massive star forming regions are, in general, much more complex
than in regions which form only low mass stars. For instance, in the early stages within
a low mass star forming region, the dynamics can be understood in terms of a few broad
categories: large scale infall, which causes a disk to form, accretion through the disk, and
outflow to release angular momentum (see for example, Di Francesco et al. 2001, André
et al. 1993, Muzerolle et al. 2003). In intermediate and high mass star forming regions,
http://arxiv.org/abs/0704.1245v1
– 2 –
turbulence, stellar winds, multiple sites of star formation, and, for regions with massive star
formation, the presence of HII regions, all contribute to the dynamics in these regions as
well (i.e. Beuther et al. 2006, Shepherd & Churchwell 1996, McKee & Tan 2003, Krumholz
et al. 2005). The more complicated source dynamics make the processes involved in the
formation of the most massive stars much more difficult to understand than those involved
in the formation of lower mass stars. Adding to the complexity, massive stars do not form as
often as their lower mass counterparts and so we must look to larger distances before finding
examples of high mass star formation. For instance, the average distance to the 63 sources
in Shirley et al. (2003, hereafter S03) is 5.3 kpc.
If we assume massive stars form through accretion, that this accretion occurs in the
inner regions of disks (i.e. Pudritz & Norman 1986, or Shu et al. 1994), and that these
disks have radii of a few thousand AU (Chini et al. 2004, Beltran et al. 2004, Cesaroni et
al. 2005), we do not yet quite have the resolving power to detect accretion directly at the
distances to massive star forming regions (1000 AU at 5.3 kpc is ∼ 0.2′′). Here, we define
accretion as the infall motions from the disk onto the forming star, in contrast to the larger
scale motions of envelope material falling onto the disk. However, while we cannot observe
accretion directly, its presence can be inferred from the presence of accretion tracers such
as larger scale infall and outflow. Infall can act to replenish disk material as mass accretes
onto a protostar (Nakamura 2000), while molecular outflows serve as a release mechanism for
the angular momentum which builds up during the accretion process (e.g Arce et al. 2006).
These large scale motions are seen in star forming regions of all mass scales (see for instance
Beuther & Shepherd 2005).
In massive star forming regions, the accretion rates are orders of magnitude greater than
in low mass star forming regions (i.e. Beuther et al. 2002), while the accreted masses are
only approximately one order of magnitude greater. These accretion rates and masses result
in accretion timescales that are much shorter than in low mass star forming regions, which
allows the Kelvin-Helmholtz timescale to become important in the evolution of the protostar
(e.g., there is no pre-main sequence stage for massive star formation). The outward radiation
and thermal pressure from the forming star becomes strong enough that it can ionize the
surrounding medium and a small, highly ionized HII region (either hypercompact (HCHII) or
ultracompact (UCHII) region, Keto 2003) can form. It is still unclear whether the outward
pressure needed to create the HII region is strong enough to halt accretion, or whether
accretion can continue in some form (either through a molecular or ionized disk, or through
an ionized accretion flow) after the formation of an HII region. Some models suggest that
accretion must halt before the onset of a visible UCHII region (i.e. Garay & Lizano 1999,
Yorke 2002), while other models suggest that an ionized accretion flow can continue through
an HII region (i.e. Keto 2003, 2006).
– 3 –
There is now also observational evidence which suggests that, once the protostar be-
comes hot enough to ionize its surroundings, both modes of massive star formation (halted
and ionized accretion) are possible. G10.6-0.4 has been shown to have an ionized accretion
flow by Sollins et al. (2005) and Keto & Wood (2006), while accretion in G5.89-0.39 seems
to have halted at the onset of the UCHII region (Klaassen et al. 2006). Although sample
statistics at this point are still quite small, these two examples pose interesting questions.
We do not yet have enough data to determine whether the apparently conflicting models of
halted and ionized accretion can both be correct. However, we can begin with a uniform
survey of infall and outflow tracers in massive star forming regions in order to constrain
massive star formation scenarios.
In this paper, we present a survey of 23 massive star forming regions. Because we
are interested in the relationship between accretion and outflow after the formation of an
HII region, our source selection criteria include (1) the presence of an UCHII region, which
indicates that there is a massive protostar forming, and (2) previous evidence of outflows,
which suggests ongoing accretion in most formation scenarios. Sources were selected based
on inclusion in the Wood & Churchwell (1989) and Kurtz et al. (1994) catalogs of UCHII
regions as well as having molecular outflow signatures in the Plume et al. (1992) survey of
massive star forming regions. Additional sources were taken from Hunter (1997) which were
shown to have both UCHII regions and molecular outflows.
We describe the observations collected for this survey in Section 2, we discuss the results
of these observations in Section 3, and present our conclusions in Section 4.
2. Observations
Observations of SiO (J=8-7), HCO+, and H13CO+ (J=4-3) were taken at the James
Clerk Maxwell Telescope (JCMT)1 in 2005 (as parts of projects M05AC11 and M05BC04).
SiO (347.330 GHz) and H13CO+ (346.999 GHz) were observed simultaneously in the same
sideband by tuning the receiver to 347.165 GHz. Thirteen or twenty minute observations,
depending on the source elevation, were taken towards each source with a velocity resolution
of 1.08 km s−1, which resulted in rms noise levels of TMB <0.07 K. Separately, we observed
HCO+ (356.370 GHz) with a velocity resolution of 0.53 km s−1 to an rms noise limit of
TMB <0.13 K in twenty minute integrations. Both sets of observations were taken in position
1The James Clerk Maxwell Telescope is operated by The Joint Astronomy Centre on behalf of the Particle
Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organisation for Scientific
Research, and the National Research Council of Canada.
– 4 –
switching mode with dual mixers and a sideband rejection filter in place. Table 1 shows the
positions, rms noise limits of both tunings, the local standard of rest velocity, and distances
to all sources in this survey. The half power beam width for these observations is 15′′, and
the main beam efficiency is ηmb=0.62. Data were obtained using the DAS autocorrelator
system and reduced using the SPECX software package.
Linear baselines were removed from all spectra except for those towards G10.47. For
this source, there were so many different chemical species in the observed spectrum that we
were unable to fit a linear baseline over the entire ∼ 700 km s−1 bandwidth of the 347 GHz
observations or the ∼ 450 km s−1 bandwidth of the 356 GHz observations. In this case, no
baseline was removed.
We also present 9 × 9 maps of one source (G45.07) in the same emission lines. These
raster maps are sampled every 5” and have rms noise limits of 0.10 and 0.14 K (TMB) for
the 347 and 356 GHz tunings, respectively. Note that these values are different than the
ones reported in Table 1 for the single pointing observations. The DAS autocorrelator was
configured with the same tunings as were described above for the single pointing observations,
and the maps were centered at the same position.
3. Results
This survey of single pointing observations towards 23 massive star forming regions is
meant as an initial, uniform survey from which to base future observations. With these
observations, we can only comment on the molecular gas component within our beams; we
cannot discuss the larger scale molecular dynamics, or the ionized gas components of these
regions. For our sources, the selection criteria of having an HII region confirms that these
regions are forming massive stars.
Figures 1 through 4 show the single pointing SiO and HCO+/H13CO+ spectra towards
all sources; the spectra are ordered according to SiO integrated intensity. For each panel in
these figures, line brightnesses have been corrected for the JCMT main beam efficiency and
centered on the VLSR of the source (Table 1). Figures 1 and 2 show SiO and HCO
+/H13CO+
spectra, respectively, towards the sources with no SiO detections. Figures 3 and 4 show the
SiO and HCO+/H13CO+ spectra, respectively, towards sources with SiO detections. Peak
line strengths and integrated intensities are given for all three lines in Table 2.
SiO was only detected in 14 out of our 23 sources, where we define a detection as a
minimum of 4 σ in integrated intensity. The rms noise limits in integrated intensity were
calculated using ∆I = Trms∆v
Nchan where Trms is the rms noise level in K, ∆v is the velocity
– 5 –
resolution of the observations, and Nchan is the number of channels over which the integrated
intensity is calculated. HCO+ was detected in all sources and H13CO+ was detected in all but
two sources. Along with HCO+, we serendipitously observed SO2 (J=104,6-103,7 at 356.755
GHz) in 17 of our sources. The HCO+ and SO2 lines are only separated by 17 km s
−1 and
thus the lines were blended in eight sources. For the sources with SO2 detections, we have
also plotted (in gray) the lower intensity HCO+ observations in order to highlight the SO2
emission (Figures 2 and 4).
Double peaked HCO+ line profiles were observed towards 10 sources, with nine of them
having stronger blue peaks than red. This blue line asymmetry in an optically thick tracer
such as HCO+ is often suggestive of infall (i.e Myers et al. 1996). We discuss the possibility
of our observations tracing large scale infall further in Section 3.2.
The distances to our sources, as taken from the literature, are shown in Table 1. The
average distance is 5.7 ± 3.8 kpc, where the error quoted reflects the 1 σ dispersion in the
distances. Since our observations were taken with a 15′′ beam, this resolution corresponds
an average linear size of 0.4 pc for our observations.
SiO is a well known outflow tracer, since in the general interstellar medium, Si is frozen
out onto dust grains. When the gas in a region is shocked (i.e. the gas through which a
protostellar outflow is passing) the dust grains can sublimate and Si is released into the gas
phase. After the passage of a shock, the SiO abundance ([SiO]/[H2]) can jump to almost
10−6, whereas the dark cloud abundance of SiO is often closer to 10−12 (see for example
Schilke et al. 1997, Caselli et al. 1997, or van Dishoeck & Blake, 1998).
While SiO is easily identified as an outflow tracer, the emitting region for HCO+ is much
less certain. Many authors suggest that HCO+ can be used to trace the envelope material
surrounding a protostellar region (i.e. Hogerheijde et al. 1997, Rawlings et al. 2004), while
others suggest that it traces disk material (i.e. Dutrey et al. 1997). One thing that is
apparent, however, is that it becomes optically thick very quickly and readily self absorbs.
We detected SO2 in 17 of our sources, suggesting that our beam contains at least some
molecular gas at temperatures greater than 100 K (see for instance Doty et al. 2002, Charnley
1997). However, Fontani et al. (2002) have determined the average temperature in twelve
massive star forming regions to be 44 K, using observations with beam sizes comparable to
those presented here. For seven of our sources, which were observed in the Fontani et al
(2002) sample, the average temperature is also 44 K. Thus, in the following analysis, we
adopt an ambient temperature of 44 K for all sources.
Table 3 shows the column densities for each region derived for both SiO and H13CO+.
The column density was calculated assuming that each tracer is optically thin, in local
– 6 –
thermodynamic equilibrium, and at an ambient temperature of 44 K. For optically thin lines,
the column density of the observed transition scales directly with the integrated intensity of
the line (see for instance, Tielens 2005):
8kπν2
TMBdv (1)
where Nu is the column density in the upper state of the transition, Auℓ is the Einstein A
coefficient, ν is the frequency of the J=u-ℓ transition, and
TMBdv is the integrated intensity
of the line. The column density of this one state can then be related to the total column
density of that molecule through the partition function. It is the total column density for
the molecule (not the observed state) that is presented in Table 3.
3.1. Source Properties derived from SiO observations
For each source we determined the column density, or upper limit to the column density,
in SiO (Table 3) using the methods described above. These column densities can be compared
to the column densities of other molecules (i.e. CS) for the same regions in order to determine
the fractional abundance of SiO, if the abundance of the other molecule is known. We were
able to obtain the CS or C34S column densities for fifteen of our sources from the literature.
Column densities for fourteen sources were taken from Plume et al. (1997, hereafter P97),
with the column density for one additional source taken from Wang et al. (1993). For
those sources with C34S column densities instead of CS column densities, we assumed an
abundance ratio of [CS]/[C34S] = 22 (Wilson & Rood, 1994) to determine a CS column
density. The abundance of CS, relative to H2, was calculated by S03 for 13 of these sources,
and we assume a CS abundance of 1.2 × 10−9 for the other two source for which the CS
column density is known, since this was the average CS abundance as calculated by S03.
We then compare the column density and abundance of CS to our observed SiO column
density, or column density upper limit, to determine the abundance of SiO relative to H2 in
our sources (Table 3).
Despite our source selection criteria requiring previous evidence of outflows, we detected
SiO towards only 14 of our 23 sources. This raises a number of questions, such as: is the
observed SiO in fact tracing outflow if we do not detect it in all sources? Why do we not
detect SiO in all sources? Is the signal being beam diluted at large distances? Has the Si
evolved into other species? Below we first address whether the detected SiO can be used
as an outflow tracer, and then discuss reasons for our non-detections of SiO in nine of our
sources.
– 7 –
Si is liberated in shocks, and if these shocks are not due to the outflow, they must be due
to the photo dissociation region (PDR) surrounding the UCHII region. Evidence for diffuse
(not collimated) SiO can be seen in W75N (Shepherd, Kurtz & Testi, 2004) suggesting that
the SiO may be due to the PDR and not an outflow. Models and observations of SiO in
the PDRs around high mass star forming regions suggest moderate SiO enhancement, and
that the SiO abundance is independent of the ambient radiation field (i.e. Schilke et al.
2001). Schilke et al. (2001) find SiO column densities of ∼ 1012 cm−2 in their observed
PDRs. This is, admittedly, below our detection threshold; however, we detect average SiO
column densities of ∼ 1014 cm−2. This suggests a possibly higher SiO abundance than found
in PDRs.
The enhanced SiO column density alone is not enough to discount the origin of the SiO
in our sources as the PDR and so we can consider how our SiO abundance varies with the
ambient radiation field. For the gas near an HII region, we can approximate the strength
of the ambient radiation field using the Far Infrared (FIR) luminosity of the region. For
twelve of our sources with SiO detections and abundance calculations, we obtained the FIR
luminosity from either Wood & Churchwell (1989), Kurtz et al. (1994) or Evans et al. (1981).
We then scaled their values for the different source distances used in this study (see Table 3).
Comparing the SiO abundance to the source luminosity (see Figure 5), we find that the SiO
abundance increases with source luminosity. There is a 10% chance that this relationship
could arise from uncorrelated data. Thus, it is possible that our result is contrary to the
findings of Schilke et al (2001), and we suggest that the SiO we observe does come primarily
from outflow shocks.
We can also compare our detection rate of SiO to that of the SiO survey towards maser
sources of Harju et al (1998) who observed SiO (J=2-1) and SiO (J=3-2). For comparison
to our results, we only consider the sources in Harju et al. which are listed as UCHII
regions. Our detection rate is 61%, compared to their rate of 29%. While our detection
threshold is slightly lower (our observations have rms noise levels generally below 0.05 K
at 347 GHz, while their rms noise levels are generally below 0.08 K), we suggest that the
different detection rates are due to differences in the source selection criteria. Although both
samples contain UCHII regions, our sample contains sources with previous observations of
outflows, while Harju et al. have selected sources based on previous observations of masers.
For the 12 sources which overlap between the two studies, we detect SiO towards 9 sources,
while they detect SiO towards 10. They detected SiO (J=2-1) in G31.41 while we did not
detect it in SiO (J=8-7).
Based on this comparison to SiO observations of UCHII regions not selected by outflows,
which have a lower SiO detection rate, and that our SiO abundance increases with source
– 8 –
luminosity, we suggest that the SiO, in the 14 sources in which it is detected, is being
generated in the outflow. Previous, high resolution observations of SiO also suggest that
SiO can be enhanced in the outflows from high mass stars (i.e. Beuther et al. 2004, Beuther,
Schilke & Gueth 2004) just as it is in the outflows from low mass stars.
The nine non-detections in our sample could be caused by beam dilution if these sources
are on average further away. However, if we compare the distances for sources with and
without SiO detections, we find average distances of 6.3 ± 4.4 and 4.4 ± 2.5 kpc, respectively.
Thus, the non-detections cannot be attributed to larger average source distances and so beam
dilution can play only a minimal role in the non-detections of SiO.
If these SiO non-detections are not due to distance effects, there must be some local
phenomenon which can explain why SiO is not being detected in regions known to contain
protostellar outflows. It is possible that the Si is evolving into different species and the SiO
abundance is dropping back down to dark cloud values. Pineau des Forêts et al. (1997)
suggest that a few×104 yr after the Si is liberated from dust grains and forms SiO, it can
either freeze out back onto dust grains or oxidize and form SiO2. Thus, the lack of SiO may
be due to silicon moving into other species if it was liberated more than 104 yr ago. This
interpretation implies that we did not detect SiO in some of our sources because the outflow
generating mechanism shut off more than 104 years ago, and the outflow observed in HCO+
(or in other molecules by other authors) is a remnant of previous accretion.
The kinematic ages of nine of our sources are listed in the Wu et al (2004) catalog of high
velocity outflows. Of these nine sources, five were also included in the P97 and S03 studies.
This results in five sources for which we have both the kinematic age of the outflow, and
the abundance of SiO. The relationship between outflow age and SiO abundance is shown
in Figure 6 along with the model predictions of Pineau des Forêts et al. (1997). With only
five points, it is difficult to draw conclusions about the relationship between outflow age
and SiO abundance, especially given the uncertain beam filling factor. At higher resolution,
these points would likely move upwards to higher abundances. A larger beam filling factor
would move the points for the two young outflows (G5.89 and Cep A) towards the model
predictions. As for the other three outflows, this would move them further from the model
predictions. We suggest that the outflow generating mechanism is continuing to shock these
regions, replenishing the SiO. The oldest of these sources (G192.58) shows an infall signature
in HCO+, which also suggests that the outflow is still being powered.
– 9 –
3.2. Source properties derived from HCO+ observations
Given the large average distance to our sources (5.7 kpc), our 15′′ beam subtends an
average linear distance of 0.4 pc. Thus, it is quite likely that the HCO+ emitting region does
not entirely fill our single JCMT beam. In general we can determine the beam filling factor,
f , for each source using
TL = fsourceTs(1− e−τ ) (2)
where TL is the line brightness temperature measured at the telescope (corrected for telescope
efficiencies) and Ts is an approximation to the ambient temperature (44 K) which is valid
at densities greater than ∼ 103 cm−3 (Rholfs & Wilson, 1994). In a number of our sources,
the HCO+ is asymmetric, and thus cannot be consistently used to determine a beam filling
factor. Instead, we can use the optically thin H13CO+ line, and we can simplify the above
equation to f = TL/(Tsτ). These values are shown in Table 3.
HCO+ becomes optically thick quite quickly due to its relatively high abundance with
respect to H2 ([HCO
+]/[H2] ∼ 10−8), and as such, can be used to roughly trace outflow and
to trace infall (i.e. Myers et al. 1996) if the line profile shows a double peak. We determined
the optical depth of HCO+ towards each source using Equation 1 of Choi et al. (1993), and
found that in all but one case (G139.9), it is optically thick. In all cases, the optical depth
of HCO+ is less than 77 (the abundance ratio between HCO+ and H13CO+, Wilson & Rood,
1994), resulting in optically thin H13CO+ towards all sources. Because HCO+ is optically
thick in its line center, the line wings can be used to detect outflows, and so, if there is
an outflow, it should be detectable in HCO+ even if it goes undetected in SiO. Gaussian
profiles were fit to our HCO+ spectra (either single or double Gaussians, depending on the
observed line shape), and the fits were subsequently subtracted from the spectra to leave
only the residual outflowing gas. When using two Gaussians to fit the self absorbed spectra,
we employed a method similar to the single Gaussian fitting of Purcell et al. (2006) because
we used the sides of the detected lines to fit our profiles (see their Figure 3). Comparing the
two Gaussian fits to single Gaussian fits showed no significant differences in distinguishing
line wing intensity. Because of the possibility of contamination from SO2 emission in the
blue shifted outflow wing (at -17 km s−1), the peak brightness of the residual emission was
determined using only the red shifted wing emission. In all cases (except for G10.47 for
which we could not find a linear baseline), we found a minimum of a 5σ peak brightness
temperature in the residual line wing emission, with 19 of our sources having a minimum of
a 10σ peak. This result suggests that we can detect outflow motions in all sources using our
detections of HCO+, despite not detecting SiO towards every source.
– 10 –
Our observations show that for ten of our sources, the spectral line profile of the HCO+
emission has a double peak. This profile could either be due to self absorption of the optically
thick HCO+ line or from multiple velocity components within our 15′′ beam. To break this
degeneracy, we observed the optically thin H13CO+. If the H13CO+ line has a single peak at
the same velocity as the HCO+ absorption feature, then it is likely that the HCO+ line is self
absorbed. If, however, the H13CO+ also has two peaks, and they are at approximately the
same velocities as the two HCO+ peaks, it would suggest that there are multiple components
within the beam. Of our ten sources with double HCO+ peaks, only one shows a double
peak profile in H13CO+ (G20.08). This results in nine sources with double peaked optically
thick HCO+.
In addition to the nine optically thick sources, similar line asymmetries appear in a
number of other sources. However, in these sources, there is no clear emission gap producing
a double peak profile, only an emission shoulder (i.e. De Vries & Myers 2005). If we take
G75.78 as an example, the HCO+ line peak is red shifted from the rest velocity of the source,
with a blue shifted emission shoulder.
There are a number of different kinds of source dynamics that can lead to the double
peaked line profiles seen in our spectra, such as infall, outflow and even rotation. However,
infall is the only one of these processes which would produce line asymmetries which are
consistently blue (i.e. the blue peak is higher than the red peak or shoulder). If these
profiles were due to outflow or rotation, there would be no statistical reason to have more
sources with higher blue peaks than red peaks. Many previous studies have investigated the
statistical significance of using this type of optically thick blue line asymmetry to trace infall
as opposed to other dynamical motions (i.e. Mardones et al. 1997 and Gregersen et al. 1997
for low mass star forming regions, and Fuller et al. 2005 for high mass star forming regions).
Of the 10 sources in our survey which have double peaked HCO+ profiles, we suggest
eight may be indicative of infall. The other two sources are G20.08 and G45.47. G20.08
has already been shown to have multiple components in the beam from the double peaked
H13CO+ profile, and G45.47 has a brighter red peak than blue. There are two additional
sources (G19.61, and G240.3) in which HCO+ has a strong red shifted shoulder, which we
suggest may also be tracing infall. This analysis gives a total of ten infall candidates in our
sample of 23 sources.
A recent survey of HCO+ (J=1-0) towards sources with methanol masers shows an even
distribution of sources with blue and red line asymmetries, and a higher percentage of self
absorbed lines than in our study (Purcell et al. 2006). Of the six sources which overlap
between our survey and that of Purcell et al, all six are self absorbed in HCO+ (J=1-0).
Five of them have blue line asymmetries consistent with infall, while only one (G31.41) has
– 11 –
its red peak brighter than its blue peak. We only find self-absorption in HCO+ (J=4-3) for
three of these six sources. In two of the sources for which we do not see a clear self-absorption
feature, we do see evidence for a red shifted shoulder which may be showing unresolved infall.
The sources in Purcell et al. (2006) have an even distribution of red and blue line
asymmetries, while we have a clear bias towards detecting blue line asymmetries. This
comparison could suggest that the higher energy J=4-3 transition of HCO+ is a better
tracer of infalling gas because it does not self absorb as readily as the J=1-0 transition.
For each of the 8 sources with blue, double peaked HCO+ profiles, we can determine an
infall velocity (vin) using the two layer radiative transfer model of Myers et al. (1996). Using
their equation 9, we find infall velocities for all eight double peaked infall sources (Table 4).
The mass infall rate can then be determined using:
ρV vin
πnH2µmHr
gmvin (3)
where µ is the mean molecular weight (µ = 2.35), the geometric mean radius (rgm) is the
unresolved circular radius of the HCO+ emitting region derived from the beam radius and
the beam filling factor (rgm =
frbeam), and nH2 is the ambient source density. For seven
of the sources in Table 4, the ambient density was determined by either P97, Hofner et
al. (2000), or Wang et al. (1993). We could not find the ambient density for the eighth
source (G192.58). From this analysis, we determine mass infall rates ranging from 1×10−2
to 2×10−5 M⊙ yr−1. These values are slightly higher than those generally observed for low
mass star forming regions, but are consistent with the accretion rates derived for high mass
star forming regions by McKee & Tan (2003). Since outflow rates are orders of magnitude
higher in high mass star forming regions (i.e. Beuther et al. 2002) it is not unreasonable to
suggest that infall rates are also much higher in these regions.
The mass outflow rate for only one of these sources (Cep A) can be determined from
the Wu et al. (2004) survey of high velocity outflows by dividing the mass in the outflow
by the kinematic age of the outflow. We find the ratio of the mass outflow rate to the mass
infall rate to be Ṁout/Ṁin ≈ 16. This value is only slightly higher than values seen in other
high mass star forming regions (i.e. Behrend & Maeder 2001). Also, models suggest that a
mass equivalent to 20-30% of the mass accreted onto a protostar is ejected as a wind (i.e.
Pelletier & Pudritz 1992, Shu et al. 1994), and that this wind entrains 5-20 times its mass
in the outflow (Matzner & McKee 1999).
– 12 –
3.3. Source properties derived from mapping G45.07
At a distance of 9.7 kpc, G45.07 is one of our furthest sources. This source was known
from previous observations to have multiple continuum sources (De Buizer et al. 2003,
2005). The three continuum sources were observed in the mid-Infrared (MIR) and all three
fall within 6′′ of our map center. A fairly young outflow has also been mapped at high
resolution (< 3′′ synthesized beam) in CO and CS towards this region (Hunter et al. 1997).
They observed a bipolar outflow with a position angle of -30◦ (east of north), as well as a
red shifted absorption feature in their CS observations which they take to be indicative of
infall.
Due to the large distance to this source, we should be able to detect all of the emission
associated with this source in a fairly small map. In the left panel of Figure 7 we present a
map of the SiO (contours) and H13CO+ (halftone) emission in this region. The right panel
of Figure 7 shows the HCO+ emission. In both figures, the 5σ (2.1 K km s−1) H13CO+
emission contour is plotted as a dashed line to help guide the eye. The first contour for SiO
in the left panel and only contour of HCO+ in the right panel are also 5σ (2.2 and 3.5 K
km s−1 respectively). The differences in the 5σ contour levels for each tracer come from the
different single channel rms levels between the two tunings, and the width of each line as
given in Table 2. Also plotted in both figures are the three MIR continuum sources observed
by De Buizer et al (2003, 2005).
If we did not have the added information provided by this map, there would be two
main conclusions we could draw from our single pointing observations towards this source.
The first is that the enhanced blue emission in all three tracers suggests that our pointing
is observing more of the blue shifted outflow lobe than the red. Second, since we only have
one peak in the spectrum of each tracer, there is only one source and we cannot classify it
as infalling.
The left panel of Figure 7 shows contours of SiO emission superimposed on the H13CO+
halftone. With beam spacings of 5′′, these maps are oversampled; however we note that much
of the structure in the SiO emission is on scales comparable to the size of the JCMT beam.
For instance, the structure at ∆α = −10′′, ∆δ = 0′′ is offset from the map center by more
than the radius of our beam and could be independent from the emission at the map center.
There is also SiO emission at ∆α = 5′′, ∆δ = 15′′, which is more than a full beam away
from the map center, and suggests that the SiO emission is more extended that the primary
beam of our observations. In fact, it appears as though there is a second SiO emission peak
towards the upper left of the left panel of Figure 7. Interestingly, there does not appear to
be as much H13CO+ emission at this northern position. This comparison shows that the
SiO and H13CO+ lines are tracing different gas populations in this region. The excess SiO
– 13 –
emission is offset from the map center in the same direction as the CO emission shown in
Hunter et al. (1997) at much higher resolution.
The right panel of Figure 7 shows the HCO+ emission for this region. It appears that
the HCO+ emission extends much further than the SiO emission, suggesting it is tracing the
larger scale envelope material. The line through the middle of this plot indicates the cut
taken for the position-velocity (PV) diagram along the outflow axis as described by Hunter
et al. (1997)
The two panels of Figure 8 show the PV diagrams for SiO and HCO+ in our maps both
perpendicular and parallel to the outflow axis defined by Hunter et al. (1997). Our single
pointing HCO+ spectrum (Figure 4) suggests we are observing more blue shifted outflow
emission than red shifted emission; however, from our PV diagrams, we see that there is
excess blue emission at all positions in our map. This excess blue emission cannot be due
to outflow alone; instead, it could be due to an inherent velocity shift between the three
continuum sources in our beam. We can, in fact, fit three Gaussian components to most of
our HCO+ spectra. These Gaussians peak at velocities of 60, 52 and 44 km s−1, with the
peak temperature for each component decreasing with velocity. The third component (at 44
km s−1) could not be fit at all positions because it was intrinsically weaker than the other
two peaks, and was lost in the noise towards the edges of the map. It appears as though
this third component might be contamination from SO2, which should occur at an apparent
velocity of 41 km s−1 (or -17 km s−1 in Figure 4).
Perpendicular to the outflow axis, the mean velocity of the HCO+ line peak appears
to shift from ∼ 58 km s−1 at an offset of +15′′ from the source center to ∼ 61 km s−1 at
an offset of −15′′ from the source (Figure 8). Given our velocity resolution (1.08 km s−1)
and spatial resolution (15′′), is unclear whether this velocity shift is real. If it is, it could
indicate large scale (∼ 1.4 pc) rotation within the core, on a much larger scale than would
be expected for a rotating accretion disk.
3.4. Correlations between Datasets
Previously, we discussed the reasons why we do not detect SiO towards a number of
our sources, and have calculated the mass infall rates for the sources with double peaks
in their HCO+ emission, but we have not yet discussed the correlations between the two
species. In Figure 9 we plot the logarithm of the abundance of H13CO+ against the same
quantity for SiO (open circles), as well as the column densities of both species (filled circles).
The abundance of H13CO+ was calculated in the same manner as the abundance of SiO
– 14 –
described above (using the CS column density from P97 and the CS abundance from S03).
The probability of obtaining these correlations if the data are, in fact, uncorrelated is 6×10−3
for the abundances, and 2×10−4 for the column densities.
As stated earlier, SiO is a well known shock tracer, and as such, an increased abundance
of SiO would suggest more shocked material within our 15′′ beam. The (generally) infall
tracing HCO+ has been shown by some authors to be destroyed in strong shocks (i.e. Bergin
et al. 1998, Jørgensen et al. 2004). However, Wolfire & Königl (1993) suggest that HCO+
can be enhanced in regions with high energy shocks, where electron abundances are much
greater. This enhancement in the electron abundance increases the formation rate of ions,
and we suggest that this is responsible for the H13CO+ abundance enhancement in our
sources. This correlation between the abundances of H13CO+ and SiO suggests that HCO+
and H13CO+ are not only tracing infalling gas, but also the outflowing gas as well. This
conclusion is supported by the strong, and broad, line wing emission detected in HCO+ (See
Section 3.2).
HCO+ over abundances have been seen in high mass star forming regions not included
in this study like NGC 2071 (Girart et al. 1999) and Orion IRc2 (Vogel et al. 1984). In
these two papers, the over abundances of H13CO+ are with respect to ambient cloud tracers
such as CO and H2, rather than the high density or shock tracers like the CS and SiO with
which we are comparing our H13CO+ abundances. However, Viti & Williams (1999a,b) show
that HCO+ is indeed over abundant with respect to CS in the gas surrounding HH objects,
and Jiménez-Serra et al. (2006) also show that the abundance of H13CO+ can be enhanced
with respect to SiO by up to a factor of ten in the same regions ahead of HH objects.
4. Discussion and Conclusions
Without maps of each region, it is impossible to tell how much of the HCO+ emission in
the line center and in the line wings is due directly to infall and outflow motions; however,
based on the arguments we have presented above, we suggest that ten sources show infall
motions, and all 23 source show outflow motions based on the HCO+ line profiles. We have
found evidence for recent outflow activity (SiO emission) in 14 out of our 23 sources. Seven
of these outflow and infall sources overlap. M17S, G192.6 and G240 appear to show only
infall signatures and no SiO outflow signatures. They do, however, appear to have HCO+
outflow signatures of a minimum of 8σ.
Detection of line wing emission in HCO+ and the relationship between H13CO+ and
SiO abundances described in the previous section suggest that while SiO is tracing outflow
– 15 –
in most sources and HCO+ is tracing infall in some sources, HCO+ is also observable in the
outflowing gas for all regions.
We find that the non-detection of SiO in nine of our sources is not due to beam dilution
or larger average distances to the source, but possibly to older outflows for which the Si has
likely either frozen back onto dust grains or evolved into SiO2. For these sources, it appears
that the accretion may have ceased, and the observed outflowing gas is a remnant of previous
accretion.
We have found seven sources with SiO outflow signatures but no infall signatures in
HCO+. This result could be due to a number of factors such as beam dilution of the
infalling gas which masks the spectral line profile we would expect for large scale infall.
It is possible that, as the outflow ages and widens, it may impinge on the region in
which we could detect infalling gas. For outflow cones oriented along the line of sight, the
younger, narrower outflows would have infalling gas with large line of sight velocities and be
likely to produce an observable infall signature. However, for the older outflows which have
widened, the largest infall velocities will be in the plane of the sky, and unobservable at the
resolution of the JCMT. However, this effect would not be as pronounced for outflows in the
plane of the sky (such as G5.89, for which we do not see an infall signature). It is difficult to
assess the importance of this effect without detailed information on the outflow orientation
in each source.
Thus, we suggest that some of these sources may have finished accreting, and what we
observe are remnant outflows from a previous phase of accretion. This scenario was suggested
by Klaassen et al. (2006) to explain the large scale outflow in G5.89, and this source is one
of these seven sources with an SiO outflow and no apparent infall signature.
The seven sources which show recent outflow activity (those with SiO emission) and
which appear to be undergoing infall are suggestive of ongoing accretion beyond the onset of
the HII region. If accretion is ongoing in the presence of an HII region, then it seems likely
that this accretion flow may be ionized. This ionized accretion scenario could be similar to
low angular momentum accretion with high ionization as suggested by Keto (2006). Thus,
from this survey, we suggest that both models of ionized accretion and halted accretion may
be important in describing the evolution of a massive protostar (or protostars) beyond the
formation of an HII region.
We would like to acknowledge the support of the National Science and Engineering
Research Council of Canada (NSERC). We thanks the referee for helpful comments which
improved the paper. P.D.K would also like to thank E. Keto and D. Johnstone for helpful
– 16 –
discussions during the preparation of this manuscript.
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Table 1. Observed Sample of Massive Star Forming Regions.
Name Position (J2000) RMS noise limit (K) VLSR Distance P92 Name
RA DEC 347 GHz 356 GHz (km s−1) (kpc) ref
G5.89 18 00 30.3 -24 03 58 0.060 0.111 9 2 1 W28A2 (1)
G5.97 18 03 40.4 -24 22 44 0.044 0.069 10 2.7 5 · · ·
G8.67 18 06 19.0 -21 37 32 0.068 0.074 36 8.5 2 8.67-0.36
G10.47 18 08 38.4 -19 51 52 · · · b 0.100 67 12 2 W31 (1)
G12.21 18 12 39.7 -18 24 21 0.042 0.076 24 16.3 2 12.21-0.1
M17S 18 20 24.8 -16 11 35 0.044 0.077 20 2.3 5 M17 (2)
G19.61 18 27 38.1 -11 56 40 0.048 0.065 43 4.5 1 19.61-0.23
G20.08 18 28 10.4 -11 28 49 0.044 0.068 42 4.1 1 20.08-0.13
G29.96 18 46 03.9 -02 39 22 0.044 0.079 98 9 1 W43S
G31.41 18 47 33.0 -01 12 36 0.044 0.061 97 8.5 1 31.41+0.31
G34.26 18 53 18.5 01 14 58 0.047 0.131 58 3.7 1 W44
G45.07c 19 13 22.1 10 50 53 0.044 0.073 59 9.7 1 45.07+0.13
G45.47 19 14 25.6 11 09 26 0.037 0.063 58 8.3 6 · · ·
G61.48 19 46 49.2 25 12 48 0.037 0.048 12 2 1 S88 B
K3-50A 20 01 45.6 33 32 42 0.035 0.077 -24 8.6 3 K3-50
G75.78 20 21 44.1 37 26 40 0.038 0.068 0 5.6 1 ON 2N
Cep A 22 56 17.9 62 01 49 0.031 0.076 -10 0.7 1 CEP A
W3(OH) 02 27 03.8 61 52 25 0.026 0.079 -48 2.4 2 W3 (OH)
G138.3 03 01 29.2 60 29 12 0.066 0.073 -38 3.8 1 S201
G139.9 03 07 23.9 58 30 53 0.058 0.071 -39 4.2 1 · · ·
G192.58 06 12 53.6 17 59 27 0.037 0.053 9 2.5 3 S255/7
G192.6 06 12 53.6 18 00 26 0.074 0.098 9 2.5 3 S255/7
G240.3 07 44 51.9 -24 07 40 0.035 0.048 68 6.4 4 · · ·
aName given to source in Plume et al. (1992). When included in P97 and S03, these were the
source names used.
bToo much chemistry in spectrum to determine a reliable rms noise level.
– 21 –
cThe rms noise limits for the two maps of this source are 0.1 K and 0.14 K for the 347 GHz
and 356 GHz pointings respectively.
Note. — The source velocity here is the reference velocity of each source, which has been
removed from each spectrum. RMS noise levels are given in TMB. References:(1) Hanson et al.
2002, (2) S03, (3) Kurtz et al. 1994, (4) Kumar et al. 2003., (5) Churchwell et al. 1990, and (6)
Hofner et al. 2000
– 22 –
Table 2. Integrated Intensities for SiO, HCO+ and H13CO+.
Name SiO Properties HCO+ Properties H13CO+ Properties
TMBdv dv TMB
TMBdv dv TMB
TMBdv dv
G5.89 2.3 75.3± 0.7 120 39.2 690.0±0.8 105 8.9 65.2± 0.4 35
G5.97 <0.1 <0.3± 0.1 8 12.4 50.5 ±0.2 15 0.6 1.6 ± 0.1 8
G8.67 0.2 1.6 ± 0.3 13 13.8 80.0 ±0.3 38 5.6 20.9± 0.2 10
G10.47 1.0 9.9 ± · · · 22 7.7 79.3 ±0.4 35 1.2 9.2 ± · · · 15
G12.21 0.2 2.5 ± 0.2 25 10.0 93.8 ±0.3 30 0.8 6.5 ± 0.2 15
M17S <0.1 <0.7± 0.2 15 16.2 79.0 ±0.2 15 2.9 9.4 ± 0.2 11
G19.61 0.8 15.9± 0.3 45 10.8 162.2±0.4 55 1.8 16.1± 0.2 20
G20.08 0.4 6.9 ± 0.3 40 6.5 75.0 ±0.3 40 1.1 9.7 ± 0.2 25
G29.96 0.5 7.9 ± 0.3 40 23.6 160.4±0.3 30 3.8 13.2± 0.1 10
G31.41 <0.1 <0.7± 0.2 15 1.6 6.8 ±0.2 13 <0.1 <0.3± 0.1 3
G34.26 1.3 26.0± 0.4 55 27.8 228.2±0.6 35 4.3 33.2± 0.2 25
G45.07 0.7 10.3± 0.3 37 15.0 169.4±0.4 45 1.5 11.2± 0.3 33
G45.47 <0.1 <0.3± 0.1 12 10.1 59.4 ±0.2 20 1.3 6.1 ± 0.1 12
G61.48 <0.1 <0.3± 0.1 8 10.8 58.5 ±0.1 18 1.5 4.7 ± 0.1 8
K3-50A 0.3 2.1 ± 0.1 16 18.6 154.8±0.3 25 1.8 10.6± 0.2 18
G75.78 0.3 2.2 ± 0.1 13 12.5 136.8±0.3 45 2.5 10.3± 0.2 17
Cep A 0.5 7.3 ± 0.2 45 23.2 234.7±0.4 60 4.5 23.4± 0.2 28
W3(OH) 1.4 14.5± 0.2 50 18.5 148.2±0.3 30 2.4 11.6± 0.1 15
G138.3 <0.2 <0.3± 0.1 2 5.0 10.6 ±0.2 8 0.3 0.5 ± 0.1 4
G139.9 <0.2 <0.3± 0.1 11 8.6 18.8 ±0.1 6 <0.2 <0.3± 0.1 2
G192.58 0.3 1.3 ± 0.1 8 14.5 84.2 ±0.2 25 1.1 5.5 ± 0.1 15
G192.6 <0.2 <0.6± 0.2 30 12.0 54.1 ±0.3 15 0.8 2.2 ± 0.2 8
G240.3 <0.1 <0.4± 0.2 10 7.3 59.1 ±0.2 30 0.4 2.5 ± 0.1 14
Note. — For the sources in which we did not detect SiO or H13CO+, 3σ upper limits
on the brightness temperature and integrated intensities are given.
– 23 –
Table 3. Observed and Derived Source Parameters.
Name Detection f a Column density (×1012) [SiO]/[H2]b LFIR toutflowc
SO2 Infall H
13CO+ SiO Log(L⊙) (10
4 yr)
G5.89 Y N 0.79 39.7 421.0 -8.90 5.25d 0.2
G5.97 N N 0.28 1.0 <1.7 · · · 5.23d · · ·
G8.67 Y Y 0.24 12.7 9.0 · · · 5.70d · · ·
G10.47 Y Y 0.16 5.6 55.4 -9.53 6.26d · · ·
G12.21 Y N 0.22 4.0 14.0 -10.15 6.17d · · ·
M17S N Y 0.33 5.7 <3.9 <-10.73 5.72d · · ·
G19.61 Y Y 0.22 9.8 89.0 -9.72 5.42d · · ·
G20.08 Y N 0.13 5.9 38.6 -10.48 4.86d · · ·
G29.96 Y N 0.49 8.0 44.2 -9.12 6.30d · · ·
G31.41 N N · · · <0.2 <3.9 <-11.89 5.45d · · ·
G34.26 Y Y 0.58 20.2 145.0 -10.43 5.77d · · ·
G45.07 Y N 0.32 6.8 57.6 -9.54 6.15d 4
G45.47 Y N 0.21 3.7 <1.7 · · · 6.04d · · ·
G61.48 N N 0.23 2.9 <1.7 <-10.80 5.01d 7
K3-50A Y N 0.40 6.5 11.7 -11.60 6.35e · · ·
G75.78 Y N 0.25 6.3 12.3 · · · 5.65d 3.7
Cep A Y Y 0.47 14.2 40.8 -9.78 4.40f 0.2
W3(OH) Y Y 0.39 7.1 81.1 -11.16 5.12e · · ·
G138.3 N N 0.11 0.3 <1.7 · · · 4.57e 17
G139.9 N N · · · <0.2 <1.7 · · · 4.82e 6
G192.58 Y Y 0.32 3.4 7.3 -10.94 4.79e 50
G192.6 Y Y 0.26 1.3 <3.4 · · · · · · · · ·
G240.3 Y Y 0.16 1.5 <2.2 · · · · · · 2.3
aHCO+ beam filling factor.
bSiO abundance relative to H2 for sources with CS observations (P97) and abundance
calculations (S03).
cKinematic age of outflow from Wu et al. (2004).
– 24 –
dFar Infrared Luminosities modified from Wood & Churchwell (1989).
eFar Infrared Luminosities modified from Kurtz et al. (1994).
fFar Infrared Luminosities taken from Evans et al. (1981).
– 25 –
Table 4. Infall Velocities and Mass Infall Rates.
Name Vin
a nH2
b Ṁin
(km s−1) (105 cm−1) (10−4 M⊙ yr
G8.67 0.4 ± 0.1 1.8 4 ± 2
G10.47 1.8 ± 0.3 7.2 100± 80
M17S 1.4 ± 0.5 5.0 4 ± 2
Cep A 0.23±0.07 10.0 0.17± 0.07
W3(OH) 0.06±0.02 60.0 3 ± 1
G192.58 0.8 ± 0.3 · · · · · ·
G192.6 0.9 ± 0.4 4.0 2 ± 1
G34.26 1.5 ± 0.3 3.6 14 ± 4
aInfall velocities for sources with double peaked
HCO+ profiles.
bAmbient densities taken from P97 except for:
G34.26 (Hofner et al. 2000), M17S (Wang et al. 1993).
cMass infall rates.
– 26 –
Fig. 1.— Nine sources in which SiO was not detected (to 4 σ limits). The name of each
source is given in the top left hand corner of each panel. For each source, the source rest
velocity is plotted as ∆v = 0 km s−1.
– 27 –
Fig. 2.— HCO+ and H13CO+ sources with no SiO detections (the same sources as in Figure
1). Solid lines show HCO+ emission, while dashed lines show H13CO+ emission scaled up by
a factor of four. The temperature scale on the left hand side of the panels is the scale used
for the HCO+ spectra, and the temperature scale on the right hand side is that used for the
H13CO+ spectra. For sources with SO2 detections, we have re-plotted the HCO
+ spectra on
a larger intensity scale ( from -1 K to 3 K, in gray) to show the low lying SO2 emission. The
SO2 emission, when present, is centered at -17 km s
−1, and is indicated by an arrow. G31.41
is slightly offset to emphasize that the temperature scale has been magnified to show the
emission feature.
– 28 –
Fig. 3.— SiO spectra for 14 source in which SiO was detected (above 4 σ limits). For this
figure only, the range of temperatures plotted increases by row. The plotted temperature
ranges for each row increase by 0.5 K per row. The temperature scale in the top row extends
from 0.2 to 1 K (in the TMB scale), while the temperature scale in the bottom row extends
from 0.2 to 2.5 K. Note that the velocity scale for G5.89 is much wider (a 78 km s−1 window
as opposed to a 38 km s−1 window for the other sources) in order to show the full width of
the emission line, but that the temperature scale is the same as that for W3(OH).
– 29 –
Fig. 4.— HCO+ and H13CO+ emission from sources with SiO detections (the same sources
as in Figure 3). The solid, dashed and gray lines are the same as shown in Figure 2, as is
the placement of the temperature scales. As in Figure 3, G5.89 has been plotted separately
from the rest of the sources to stress that the velocity scale is larger for this source.
– 30 –
Fig. 5.— SiO abundance plotted as a function of source luminosity. The line of best fit shown
is given by (Log([SiO]/[H2])=0.30±0.06Log(L/L⊙)-11.7±0.3), and was calculated only for
sources detected in SiO. This shows that SiO abundance increases with source luminosity.
This is contrary to what is expected for SiO produced in PDRs, suggesting the observed SiO
is produced in outflows.
– 31 –
Fig. 6.— SiO abundance plotted as a function of outflow age. The outflow ages were taken
from Wu et al. (2004). The solid line shows the model predictions of Pineau des Foréts et al.
(1997) for the SiO abundance as a function of age (assuming Si returns to the dust grains),
and the dashed line represents the canonical dark cloud abundance of SiO. The source shown
with a downwards arrow represents G61.48, a source in which we did not detect SiO, and the
value given is an upper limit to the SiO abundance. The error bars represent 30% calibration
uncertainty between our observations and those of P97 and S03
– 32 –
Fig. 7.— Left: SiO and H13CO+ towards G45.07. The halftone scale represents the in-
tegrated H13CO+ emission, while the contours show the integrated intensity of SiO. The
velocity range used to determine the integrated intensities is the same as the single pointing
velocity range (dv), however the rms noise limits are those listed at the end of Section 2.
The first SiO contour is 5σ (3.2 K km s−1), incrementing in steps of 2σ, with the same
scale continuing into the white contours near the center. Right: HCO+ integrated intensity
towards G45.07. The solid contour represents the 5σ emission for the HCO+ emission (3.5
K km s−1), and the solid line represents the cut used for the PV diagram along the outflow
axis (PA = -30◦) presented in Figure 8. For both panels, the dashed line represents the 5σ
H13CO+ emission contour (3.0 K km s−1). The three triangles represent the Mid IR sources
detected by De Buizer et al. (2005), with the two points in the center corresponding to 230
GHz continuum sources detected at the Submillimeter Array (Klaassen et al. in prep).
– 33 –
Fig. 8.— Position-Velocity (PV) diagrams for SiO (halftone scale) and HCO+ emission
(contours) in G45.07 both perpendicular to (top panel) and along the outflow axis (bottom
panel) as defined in Hunter et al. (1997). These PV diagrams were taken at positions angles
of 60◦ and -30◦ east of north respectively, with the cut for the bottom panel of this figure
shown in the right panel of Figure 7. The first contour for HCO+ is 5σ, or 0.7 K since the
rms noise limit for this map is 0.14 K, and the contours increase in increments of 5σ. The
three dashed lines in the bottom panel show the peak velocities of the Gaussian fits to the
HCO+ spectra.
– 34 –
Fig. 9.— The abundance of H13CO+ with respect to the abundance of SiO (open
circles, dashed line of best fit) appears to increase faster than the respective column
densities of these two species (filled circles, solid line of best fit). The equations of
the two lines of best fit are: Log[X(H13CO+)]=(0.84±0.09)Log[X(SiO)]-(2.4±0.9), and
Log[N(H13CO+)]=(0.5±0.1)Log[N(SiO)]+(5±1). This relationship could be due to HCO+
being enhanced (similarly to SiO), as discussed in the text.
	Introduction
	Observations
	Results
	Source Properties derived from SiO observations
	Source properties derived from HCO+ observations
	Source properties derived from mapping G45.07
	Correlations between Datasets
	Discussion and Conclusions
