The spectra obtained in the line survey are presented in order of
increasing frequency in Fig. 4 (click here). Tables 7 (click here) to 14 (click here) summarize the resulting
fits to the line profiles, whereas Table 3 (click here) contains the beam-averaged
column densities. In the following, the results will be discussed per
molecule, describing first the data for IRS4, then for IRS5 and finally
for W 3(
).
The most recent CO maps of W 3 include those by Hasegawa et al. (1994)
in 
and 
and 6-5 lines and by
Oldham et al.
(1994) in 
. The data obtained here on additional isotopic
species at just a few positions form a complement to this study, since
mapping is beyond the scope of this work.
The CO 3-2 line toward W 3 IRS4 is strong with 
, but the CO profiles decrease sharply on the red side, which was
already noted by Mitchell et al. (1991) and HMMT. In some
of the spectra shown here the profile is further affected by emission at
the off-position, seen as an extra ``absorption''. No other molecule or
CO isotopomer shows a similar line shape. The profile can be affected
by absorption by cold foreground material, but it also bears some
resemblance to shocked profiles like those found in IC 443
(van Dishoeck
et al. 1993). In the case of IRS4, such a shock could be
generated at the interface region where the ionization front runs into
the cloud. Because IRS4 is at the back, the shock would be coming
toward us, consistent with the blue wing. HMMT also find evidence for
moving gas associated with IRS4 from pedestal features in the spectra,
which they interpret either as outflow or due to the expansion of the
HII region. The main beam temperature of 
is close to
the kinetic temperature of 55 K found from the formaldehyde lines.
Because the line is heavily optically thick, it suggests that the warm
gas fills most of the beam, and that the kinetic temperature is likely
somewhat higher than 55 K.
The CO 3-2 spectrum toward W 3 IRS5 shows strong wings, caused by the massive outflow known from other studies (e.g., Mitchell et al. 1991, 1992; Choi et al. 1993, HMMT). Another interesting feature is the strong self-reversal of the 2-1 and 3-2 profiles toward IRS5. Even in the 3-2 line, the absorption is very strong, implying a large amount of cold material in front of IRS5. Our absorption is deeper than that found in HMMT, which may be partly due to lower spectral resolution in the HMMT data. Also, a slight pointing offset from the IRS5 position results in a less strong absorption and a less symmetric profile, as is well illustrated by Figs. 1 (click here) and 2 (click here) of HMMT. HMMT estimate an excitation temperature of 26 K for the foreground gas at the IRS5 position, and find that this temperature is rather constant over the cloud. Note that this value is similar to that found from the ro-vibrational absorption line spectroscopy toward IRS5 (HMMT; Mitchell et al. 1990). The depth of the absorption in our data suggests that even colder or lower excitation material is present along the line of sight.
At least three Gaussian components are needed to fit the central line,
its wings and the self-absorption, making the fit somewhat uncertain.
It is, however, difficult to obtain any fit for the central emission
with 
smaller than 100 K, which again is close to the
kinetic temperature found from the 
lines (see Paper I),
indicating that this gas fills a significant fraction of the beam.
Additional evidence for these high temperatures stems from high-J CO
9-8, 12-11, 14-13 and 16-15 measurements (Boreiko & Betz
1991; Betz & Boreiko 1995). On the other hand, no
good fit can be obtained for 
larger than 200 K. Thus, the
hotter gas with 
visible in the
ro-vibrational CO absorption lines (Mitchell et al. 1990);
HMMT) must originate in a volume which is smaller than the 15
\
beam.
The self-reversal of the CO 3-2 line toward W 3(
) is less strong
than that toward IRS5. Several absorption components appear to be
present, although their appearance is again affected by the small
beam-switch of +180
. The 3-2 line profile also shows strong
wings. Although no map of the red and blue components has yet been
made, the movement of the water masers (Alcolea et al. 1992;
Reid et al. 1995) suggests that an outflow is present at the
position. However, the close proximity of the compact HII region W 3(OH) and its unknown contribution to the line profiles can
complicate the picture. The fit to the central component is even more
difficult and no reliable upper limit can be given. A lower limit to
the main-beam temperature is again 100 K, which is probably close to the
kinetic temperature of the warm core surrounding both W 3(OH) and
W 3(
). Higher temperatures (
) may be present, but are
not observed, because the emission from the core is already optically
thick.
The data on the CO isotopomers for the three sources are much easier to
analyze, since they do not show self-absorption and are almost Gaussian
in shape. For all three sources, even the 
lines are still
optically thick. Only the 
emission is surely optically thin.
CO column densities have been derived using the physical parameters
found in Paper I and the isotope ratios of Sect. 4.4. The 
column
densities are subsequently determined using 
, as found for the warm star-forming region NGC 2024 by Lacy et
al. (1994). The resulting beam-averaged total CO and 
column densities are summarized in Table 2 (click here). Note that the 
column
densities are a factor of 3 lower than those listed in Paper I because
of the higher adopted CO abundance.
As discussed in Sect. 3, the 
line has been mapped
interferometrically in the direction of W 3(OH)/W 3(
). The
emission is concentrated in a 
clump toward W 3(
).
However, most of the single dish emission (
) is resolved out,
suggesting that indeed the emission observed at the JCMT largely fills
the beam.
|
Source |  |  | W |  | n(H ) |
N(C O) | N(CO) | N(H ) |
| (K) | (km s ) | (K km s ) | (K) |
(cm ) | (cm ) | (cm ) | (cm ) | |
|
W 3 IRS4 | 3.73 | 3.6 | 14.3 | 55 |  |  |
 |  |
| W 3 IRS5 | 4.47 | 4.1 | 19.5 | 100 |  |  |
 |  |
W 3( ) | 3.23 | 4.4 | 15.1 | 100 |  |  |
 |  |
|
|
Using 
determined by
Lacy et al. (1994) for
NGC 2024.
The molecules in this section are all linear nitrogen-bearing molecules. Due to the non-zero nuclear spin of the nitrogen atom, these molecules show hyperfine splitting of the lower rotational lines, but in the higher lines the splitting is generally smaller than the line widths, so that these blended components are treated as a single line.
CN is one of the few molecules which is more prominent in the direction
of W 3 IRS4 than in the other two sources, as can be seen in several
spectra around 337 and 340 GHz. Some satellite lines well separated
from the main lines can be detected, but do not allow an accurate
determination of the optical depth. A complication in determining the
CN rotational excitation is that the lines in the 230 and 345 GHz
windows stem essentially from just two energy upper levels, both lower
than 40 K. The column density derived from the rotation diagram method
is therefore very uncertain. Excitation calculations have been
performed for this molecule, using a more extended set of collision
rates of Bergman (1995, private communication). The
/
ratio has been used to
restrict the gas density to 
toward IRS4,
assuming 
. Note that the 345 GHz main lines become slightly
optically thick (
).
For IRS5 the observed line ratio is the same as that found for IRS4,
implying a similar density of 
, even for the
slightly higher kinetic temperature of 100 K. The beam-averaged column
density is a factor 2 lower than in IRS4. For W 3(
), the column
density is 
, similar to IRS5.
Early interferometer measurements by Wright et al. (1995) and
single-dish observations of Dickel et al. (1980) and
Hayashi et al. (1989) have shown that there is a large
concentration of HCN south of IRS4, and a lack toward IRS5. With the
current improved receiver technology, not only HCN but also the 
and 
isotopomers have been detected toward all three sources,
implying high optical depths in the lines of the main isotope. The best
candidates for statistical equilibrium calculations are therefore the
and 3-2 lines, but unfortunately the 4-3 line is
severely blended with 
, even in the case of IRS4 where little
present. Therefore the physical parameters derived in Paper I
were used to fit the measured strength of the 
line.
The corresponding 
beam-averaged (20'') column density for
IRS4 is 
, whereas the 
isotopomer
) gives 
, resulting in
and 
,
respectively. The difference between the two numbers is well within the
error bars. The excitation temperature inferred from the statistical
equilibrium calculations is 26 K. The fact that the optically thick
main isotopomer line (J=4-3) has a main beam temperature of 5.6 K
implies a source size of more than 9.5
.
For IRS5 the same procedure was followed giving 
and 
,
which correspond to 
and 
respectively. Again, the strength of the optically thick line suggests
that the emission is beam filling. The failure of Wright et al.
(1995) to detect HCN emission in this direction may thus also
stem from the fact that the emission is not very concentrated and
therefore filtered out by the interferometer.
For W 3(
) the beam-averaged column densities are 
and 
, resulting in 
and 
. The large difference suggests that even 
emission is still somewhat optically thick. If the emission comes from
a small source, the inferred column densities are much higher and only
the 
lines stay optically thin. A small HCN source (
) is consistent with the interferometer observations of Turner &
Welch (1984). However, some of the emission must also come
from the more extended surrounding envelope, since the excitation
temperature combined with the observed main beam temperature of the
optically thick HCN indicates a source of 
.
HNC is of interest for comparison with HCN since the HNC/HCN abundance
ratio has been found to differ significantly between cold dark clouds,
where it is close to unity, and warm clouds, where HCN can be more
abundant by up to two orders of magnitude (Irvine et al. 1987;
Goldsmith et al. 1986; Schilke et al. 1992). In W 3, HCN is also
observed to be more prominent than HNC. Lines of HNC and 
were detected toward all three sources, while there is a detection of
one 
line toward W 3(
). The lines of the main
isotopomer are most likely optically thick, thereby necessitating the
use of the 
isotopomer for statistical equilibrium calculations.
The density can in principle be derived from the 4-3/3-2 ratio, but
the 
lines are weak and the signal-to-noise ratio is poor.
Also, the 3-2 line can be affected by a blend with 
, especially
toward IRS5. Nevertheless, the line ratios of order unity indicate high
densities, consistent with or perhaps even somewhat larger than those
derived in Paper I from 
.
If the physical parameters of Paper I are used to fit the 4-3 line,
toward IRS4 and thus
. The source size is found to be
. For IRS5, 
) is found to be 
and thus the beam-averaged HNC column density is
. The (uncertain) source size is
, smaller than that found for HCN. Toward W 3(
),
and thus 
. Using these parameters the inferred source size is
9
. The corresponding HCN/HNC column density ratios for the three
sources are 
, 
, and 
, respectively.
Although 
is a common molecule in the Orion and Taurus clouds, it
is still undetected in the W 3 GMC core. The non-detections toward IRS4
and IRS5 may have several reasons. The most important one is that the
lines available in the 230 and 345 GHz windows originate from levels
that are rather high up in energy (
), making it
difficult to populate these levels in gas with temperatures below 100 K.
Moreover, the molecule has a large dipole moment, so that the critical
densities are large (
). The upper limits on the lines
still provide useful limits on the beam-averaged column density, if the
temperature and density from Paper I are adopted. For IRS4 the best
upper limit comes from the 24-23 line and is 
. For IRS5 the same line gives 
. The 
column densities are an order of
magnitude less than those of HCN toward the two sources.
is detected toward W 3(
), and is one of the few more
complex unsaturated molecules identified here. Although one of the
detected lines is blended, it still gives an upper limit on density (
) and on temperature (
). These limits are
not very stringent, because an error of 30% can already bring these
values down to the physical parameters of Paper I. Using the Paper I
values gives a beam-averaged column density of 
.
HNCO is not detected toward IRS4, but is marginally present toward IRS5
and clearly toward W 3(
). Care has to be taken in analyzing its
excitation, since it is well known from observations of the Galactic
Center that infrared pumping can affect the populations (Churchwell et
al. 1986). Since only a few lines have been detected in
this survey, a study like that of Churchwell et al. is not possible,
even though intense far-infrared radiation is present toward these two
sources and pumping is likely. An alternative is the rotation diagram
method, which gives 
and 
toward W 3(
).
Toward IRS5, the rotation temperature was assumed equal to the kinetic
temperature from Paper I, because only two weak lines have been
observed. The fit of the strength of the 351.633 GHz line results in a
(highly uncertain) beam-averaged column density of 
. If 
is assumed for IRS4,
the non-detections give an upper limit 
.
CCH is a linear molecule with strong doublet lines toward IRS4. The
measurements allow determination of the density from the 4-3/3-2 line
ratio, and the results are consistent with those found in Paper I.
Using these values, the beam-averaged column density is 
. Toward IRS5, the molecule is also detected and the
line ratio indicates a similar density as found from formaldehyde. The
corresponding column density is 
.
Toward W 3(
), the 349 GHz lines of CCH are blended with the
lines, but the 
contribution can be
estimated rather well so that this does not pose a problem but only
increases the error bars. The density derived from the line ratio is
again consistent with that found in Paper I and implies 
.
is a non-saturated molecule, which can be made quite easily
through pure gas-phase chemistry. Three lines have been detected toward
IRS4, but the corresponding rotation temperature is uncertain since one
of the lines is on the edge of the spectrum. The results are 
and 
with an uncertainty of a factor of two (see Fig. 2 (click here)). 
is
not detected toward IRS5, but an upper limit of 
is found if an excitation temperature of 50 K is assumed.
Many lines are detected toward W 3(
), allowing a more accurate
rotation diagram than for IRS4. The inferred parameters are 
and 
. The rotational temperature seems to indicate that
the 
resides in the same gas as the CO (
), and
not in the compact ``hot core'' with a much higher temperature, since no
lines were detected from very high energy levels (e.g. like methanol,
see Paper I).
Formaldehyde and methanol are two particularly useful organic molecules,
which can provide a large amount of information about the physical and
chemical characteristics of the gas they reside in. Other (saturated)
organics like 
, 
and 
are typical for
``hot-core'' type regions and are only detected toward W 3(
).
Ethanol (
) also belongs in this category but has not been
detected in this survey.
The results of Paper I for IRS4 and IRS5 have not changed, but more
information on the 
isotopomer has been obtained toward
W 3(
). The 
line has been detected in two
independent sets of observations and indicates that the
ratio of integrated intensities lies in between 10
and 40. The column density for ortho-
is 
, which implies 
for an ortho/para ratio of 3. This value is consistent with
derived in Paper I.
The source size determined from the (slightly) optically thick 
line is larger than the 220 GHz beam, suggesting that
the emission comes largely from the warm core surrounding W 3(
)
and W 3(OH). Some 
is likely to be present as well in the
compact ``hot core'', but observations of higher excitation lines and/or
isotopomer are needed to probe this region.
Methanol was discussed extensively in Paper I, and the results remain
valid except for small corrections based on more extensive data. See
Table 3 (click here) for beam-averaged column densities and rotational temperatures.
It should be noted that the values for IRS5 are still uncertain, since
the detections have large error bars. The methanol lines have been
grouped in Tables 8 (click here)-10 (click here). For W 3(
) the lines are grouped according
to the quantum number of the torsional mode (
). The
different sets were fitted separately, but this did not improve the fit,
nor was any systematic trend visible. Therefore, just as in Paper I, a
single temperature was used and the scatter in the diagram is explained
best with the methanol being subthermally excited.
The 
lines toward W 3(
) were fitted with 
and 
with
large uncertainties. Especially the two lines coming from levels with
energies higher than 300 K seem to deviate from the fit. If these two
lines are left out, the rotational temperature becomes 
and
the column density 
. Thus, the column
density is fairly robust, whereas the rotational temperature is not.
The first estimate is favoured since there is no good reason to delete
the two detected lines. The scatter can come from subthermal excitation
such as found for methanol, but perhaps infrared pumping and optical
depth effects can play a rôle as well.
Statistical equilibrium calculations have been performed and show that
the kinetic temperature must be higher than 120 K and the density of
order 
. The ortho and para column densities are
1.0 and 
respectively. The ortho-para ratio
is close to unity, as expected for a warm region.
The above analysis assumed that the 
emission fills the 230 and
345 GHz beams. However, the interferometer data of Wink et al.
(1994) show that the 
emission comes from an unresolved
region, with most of the single dish flux recovered in the
interferometer. Assuming a source size of 1
, they derive a
column density 
and an excitation temperature of
105 K. If the size is indeed as small as 
, our inferred
column density would increase to 
if the
lines remain optically thin. For such high column densities, however,
optical depth effects start to become important and the column density
may well be an order of magnitude larger. Recent 
interferometer observations by Wyrowsky & Walmsley (1996, private
communication) suggest an 
column density a factor 
larger
than that listed in Table 2 (click here) in a 
region. The corresponding
abundance is 
compared with the beam-averaged
value of 
found here.
Dimethyl ether, 
, was detected toward W 3(
) through a
number of lines, which is however only a small subset of the transitions
available to this molecule. Since no collisional rates are available, a
rotation diagram was constructed (Fig. 2 (click here)). This resulted in a
rotational temperature of 
and a beam-averaged column
density of 
. Just as for 
, two
outlying points at high energies affect the fit. Although these appear
to be clear detections, a fit has also been made without them, giving
and 
.
Figure 2: Rotation diagram for 
(upper panel) toward
W 3 IRS4 (solid symbols, dashed line) and W 3(
) (open symbols,
solid line). The panels for 
, 
and 
are
for W 3(
) only. Squares represent detected lines, while triangles
denote upper limits. For an explanation of the stars and the two fits
in the 
and 
rotation diagrams, see the text
The same approach was used for methyl formate, and gives 
and 
.
Note that these column densities are only about one order of magnitude
less than found for methanol. They could be increased considerably if
most of the emission originates from a source size as small as that of
.
The chemistry of sulfur-bearing molecules, especially 
, is still
poorly understood. In this work, a large number of species containing
sulfur have been detected, especially toward IRS5.
Carbon monosulfide is well known as a tracer of the densest parts of
molecular clouds and indeed CS and 3 of its isotopomers are found toward
the three sources in several lines. Tieftrunk et al.
(1995) show in their 
of the W 3 core that
the emission is strongest toward IRS4 and southward, but that there is
relatively little J=3-2 emission in the direction of IRS5 and the rest
of the core. This trend is also found in the higher excitation lines.
Although blending is seldom a problem toward IRS4, it happens that the
line is blended with a (weak) methanol line. The
integrated line strength can, however, be used as an upper limit.
Together with the measured 7-6 line, the ratio has been employed to
constrain the density to 
, similar to that
indicated by 
. A confirmation of this value comes from the
line and the upper limit on its 7-6 line. From the
absolute strengths of the lines, the column densities are determined and
listed in Table 3 (click here). Using the cosmic abundance ratios, 
. The different isotopomers give values which agree
within a factor two, indicating that only CS is optically thick. Note
that this column density is in excellent agreement with the results of
Tieftrunk et al. (1995). The source size was estimated
from the optically thick CS 7-6 line and was found to be beam-filling,
in agreement with the map of Tieftrunk et al. (1995).
The density found from the CS isotopes toward IRS5 is also in excellent
agreement with the 
results. The column densities were derived
using a kinetic temperature of 100 K, and are listed in Table 3 (click here). From
the rare isotopomers, 
, somewhat lower
than toward IRS4. The emission was again found to fill the beam. The
column density is accurate to better than a factor of two, but it is six
times lower than the value given by Tieftrunk et al. (1995). The
reason is unclear since our column densities can explain their
measurements of the 
and 3-2 lines well for the inferred
physical parameters.
For W 3(
) the density found from the 7-6/5-4 lines is also in
excellent agreement with the formaldehyde results of Paper I. The
column densities for the different isotopomers are listed in Table 3 (click here),
and give 
. The source size was found to
be approximately 10
.
Using the 
column densities listed in Table 2 (click here), our CS abundances
are larger by up to an order of magnitude than the values given by
Tieftrunk et al. (1995) for IRS4 and IRS5, and Wilson et al. (1991)
for W 3(
).
Several isotopomers of SO have been detected toward our sources,
including 
, 
and 
. Like CS, the SO itself is
optically thick, as is also implied by the large line width of the main
lines, although outflowing gas could play a rôle. Although there are
many lines available to construct a rotation diagram, it proved
difficult to do so: the SO lines suffer from optical depth effects,
whereas the 
lines often turn out to be blended.
The excitation of SO can be studied through statistical equilibrium
calculations using the recent collisional cross sections of Green
(1994). For IRS4, the observed SO line ratios are not very sensitive to
density and temperature, but the values agree with the parameters
derived from 
, although somewhat higher temperature and density
are favoured. The beam-averaged column density is 
, when a temperature of 80 K and a density of 
are used.
Toward IRS5 not only 
and 
lines have been detected
but also some 
, indicating that the main SO lines are very
optically thick, 
. For example, the ratio of the 
lines of SO and 
is only 50 instead of the cosmic ratio
. Even the 
lines are still somewhat
optically thick. Statistical equilibrium calculations were performed
for the 
lines, resulting in 
and
. Assuming that the SO, just like
(see Paper I and below), originates in hotter (
) and
denser (
) gas than 
, the 
column
density was calculated to be 
, implying
. From the excitation calculations, the
source size is estimated to be 
. Again, this number is
somewhat uncertain, but it shows that the SO likely comes from a
compact, warm and dense region.
The SO emission from W 3(
) is very intense, just as that from
IRS5, but no 
lines are detected and only one 
line.
The statistical equilibrium calculations were therefore performed for
. The density and temperature can be constrained to 
and 
. Using the
temperature and density derived from 
, 
is found, resulting in 
. This beam-averaged column density is fairly robust
since it does not change by more than 10% if the density is increased
to 
. The source size was calculated to be large,
.
is one of the most interesting molecules found in the W 3 region.
In Paper I, its abundance was found to differ greatly between the three
sources, with IRS5 showing the strongest 
emission. Indeed, the
abundance has been found to be highly variable in other sources
as well (Groesbeck 1994; McMullin et al. 1994). In some cases such as
Orion/KL and Sgr-B2(N), the 
lines are so overwhelmingly strong
that it can be one of the major coolants in the 345 GHz range. The
excitation temperature was also found to be surprisingly high
toward IRS5. With the extended data set presented here, a high
temperature has been found toward IRS4 as well. Another big improvement
compared with Paper I stems from the recent completion of the catalogue
of 
lines.
Figure 3:
Rotation diagram for 
(open symbols) and 
(filled
symbols) in the three sources. Squares indicate lines that have been
detected; triangles denote upper limits; stars are detected lines
with 
, see text. The solid lines represent the
least-squares fits using all detected transitions. The dashed line
indicates the fit through the 
lines only. The dotted
lines represent the least-squares fits using all detected
transitions. The rotation temperatures are those derived
from all detected 
transitions (IRS4 and W 3(
)) or from the
transitions only (IRS5)
Toward IRS4, several additional 
lines have been detected, mainly
because of the better DAS backend. The inclusion of these lines in the
rotation diagram (Fig. 3 (click here)) results in 
and
. No lines from
have been detected. The column density is, within the
errors, consistent with Paper I, but the rotational temperature is
increased. The higher temperature is no longer consistent with the
kinetic temperature derived from the formaldehyde lines. This could, as
for IRS5, be due to optical depth effects, but the upper limits on the
column density are not stringent enough to provide clues
to the optical depth since 
is >3.4. More likely,
the 
emission comes from a somewhat warmer region in the direction
of IRS4. This possibility is discussed further in Sect. 6.
For IRS5, the complete data set gives 
and
, which is within the
errors consistent with Paper I. For the optically thin 
the following parameters are found: 
and
. From these
values it is clear that the ratio of the beam-averaged column densities
is not equal to the solar abundance ratio of 22. This was already found
in Paper I, where it was attributed to optical depth effects. Since the
lines have the lowest intrinsic line strengths and thus
the lowest optical depth, a separate rotation diagram has been made for
these lines. The results are 
and
. This rotational
temperature matches very well with the 
results, and the
beam-averaged column densities are consistent if the cosmic
[
] ratio of 22 is used. The fact that the rotational
temperatures are somewhat lower than found in Paper I is thus explained
by the high optical depths which flatten the slope of the fit.
In Paper I it was also argued that the lower rotational temperature of
may be caused by lack of data at higher energies. Since
then, the catalog has been completed, and this lower temperature still
persists. At the same time, our list of detected 
lines has increased and this more complete data set now nicely confirms
the 
results, giving more weight to the optical depth
argument. The 
rotational temperature of 
is no
longer consistent with the rotational temperature of 
of the
outflowing CO (Mitchell et al. 1991). However, a recent
re-analysis of the CO excitation by HMMT shows that there may be
multiple components with 
. Furthermore a kinetic
temperature of more than 200 K can still result in a rotational
temperature of 
if the 
is subthermally excited. Since
the densities toward W 3 appear to be lower than those toward the Orion
Plateau region and since 
has a large number of (optically thin)
lines to radiate through, subthermal excitation cannot be excluded. The
optically thick lines lead to a crude determination of the source size.
For 
, the most optically thick line gives a source
size of 
if the optical depth were very high. However,
the 
ratio suggests that the optical depth is still
less than unity, so that a larger source size of 
, similar
to that found for SO, is more likely.
The results for W 3(
) from the rotation diagram are 
and 
.
The 
lines give 
and 
respectively. Only few lines with 
have been identified at low signal to noise, but the same trend was
found as in IRS5, implying that the main lines are somewhat optically
thick. The inferred source size is 
.
, like 
, is a near-prolate rotor with transitions that can
be used to determine both density and kinetic temperature (see Fig. 9 of
Blake et al. 1994). Toward IRS4, two ortho and two para lines have
been detected, and good upper limits on several other lines have been
obtained. The corresponding temperature range is 
and density interval 
. Using 80 K
and 
as the best parameters, 
and 
, so
that the ratio differs somewhat from the theoretical high-temperature
ortho/para ratio of 3. The total column density 
is the same as found by the rotation diagram method.
The rotational temperature is 
, indicating
subthermal excitation.
Toward IRS5 only the 
line has been detected, prohibiting
an analysis as given for IRS4. Using the physical parameters from
, 
is found. Assuming
an ortho/para ratio of 3, this results in 
.
Many more lines were detected toward W 3(
) allowing a rotation
diagram to be constructed. The fitting parameters are 
and 
.
Statistical equilibrium calculations were performed as well, resulting
in 
and 
. The column densities are 
and 
. Again, the
ortho/para ratio differs from three but the sum of the two column
densities is close to the rotation diagram value. Most significant
however is the low temperature found from the statistical equilibrium
calculations. It suggests that the 
emission does not come from
the compact ``hot core'', but from the warmer core surrounding the
W 3(OH)/W 3(
) clumps. This interpretation is strengthened by the
fact that the emission from IRS4 comes from a similar region, which is
not compact and hot, but extended and warm.
has only few lines available in atmospheric windows, and is
therefore not much studied. Therefore the analysis of this molecule is
largely based on the single 
line at 216 GHz. Since
collisional rates for this molecule are not known and since no
rotational temperature can be derived from a single line, the excitation
temperature has to be assumed. Fortunately, the upper energy level of
this line lies at 84 K, nicely in the temperature range of the three
sources. Moreover, for excitation temperatures between 50 and 250 K,
the 
column density changes by only a factor of three. Taking
, we obtain 
; 
; and 
for IRS4, IRS5 and
W 3(
) respectively. The upper limits on the 
line at
214.7 GHz indicate that the 216 GHz line is not very optically thick
toward IRS5 and W 3(
).
For IRS5 the analysis can be improved by using the 
line
at 168.763 GHz toward this source observed by Minh et al. (1991a) in a
35'' beam. They also measured the corresponding 
line and
found the main line to be optically thick. If the two 
lines are
combined, 
and 
is derived. While the rotational temperature is of
the same order as their adopted excitation temperature, the inferred
column density is more than an order of magnitude smaller than that of
Minh et al. Using the lines of the rarer isotopomer (assuming the upper
limit at 214.7 GHz to be a two sigma detection) gives a somewhat higher
rotational temperature and a 
column density of 
. Although this value is uncertain, it suggests that
the 
column density of Minh et al. is too large. The reason for
this discrepancy is unclear. Therefore the value of
derived above is kept.
Minh et al. (1991a) also observed W 3(OH) in 
, and
their 35'' beam included W 3(
) as well. The rotation diagram
for this line and the two sigma upper limit at 214.7 GHz (used as a
detection) gives 
and 
, corresponding to 
. This number is highly uncertain, but does not
differ much from the value adopted above.
If 
toward IRS4 is lowered to 
, similar to the
values found for IRS5 and W 3(
), its column density would be
increased to 
.
OCS has not been detected toward IRS4. The best limit on the column
density comes from the 19-18 line, which gives 
for the physical parameters of Paper I.
The situation for IRS5 is somewhat better, since two lines were detected
which are reasonably far apart in energy. As can be seen from Fig. 11
of Jansen (1995), such transitions are extremely good density tracers.
The statistical equilibrium calculations give a density as high as
at temperatures of more than 200 K, but the
signal-to-noise ratio on the data (especially for the 19-18 line) is
poor. If this line were in error by a factor of two or if the source
size were small, the density would decrease to 
and
the temperature to 
. These conditions are similar to those
found for SO and 
, providing additional evidence for a separate
sulfur chemistry. If the lower temperature and density are adopted, the
beam-averaged column density is found to be 
.
Toward W 3(
) many OCS lines are detected and their strengths are
large enough to place useful limits on the 
ratio. The
non-detection of 
at the level 
indicates
that the OCS lines are not very optically thick. The OCS lines often
turn out to be blended, but in most cases OCS is the dominant component
of these blends. Therefore line ratios can again be used to determine
the physical conditions. Both the 28-27/19-18 and 29-28/19-18
ratios point at 
and 
, consistent with the 
results. The inferred
is 
.
Silicon monoxide is a special molecule, because of its intimate connection with shocks in the interstellar medium (Bachiller et al. 1991; Martın-Pintado et al. 1992). SiO has been detected toward all three sources, often in more than one line and in more than one isotopomer. Statistical equilibrium calculations of the 8-7/6-5 ratio were used to constrain the densities.
Toward IRS4, the 8-7 line is very weak. The fit of the 6-5 line and
the conditions of Paper I give a column density of 
.
Toward IRS5, the SiO emission is strong, but the 
lines are
not detected, indicating that the main isotope is not very optically
thick. A line has been found coincident with 
, but this
is probably spurious. Therefore, only the SiO lines themselves can be
used. The 8-7/6-5 ratio is large, indicating high densities. A
smaller source size of 
, comparable to that for SO and 
,
decreases both parameters somewhat, but they remain high, 
for 
. Only for much higher
temperatures would lower densities be obtained. This is not unlike the
situation close to Orion IRc2 where brightness temperatures of about
1000 K are found for the thermal SiO emission (Wright et al. 1995),
indicating a very hot and dense environment. For 
and
the inferred SiO column density is
in the 15'' beam. If the physical
parameters of Paper I are used together with 8-7 line, 
decreases to 
.
W 3(
) shows not only emission from the main isotopomer, but also
from the rarer ones. The 
lines indicate a density higher
than 
and temperatures higher than 100 K, not unlike
the 
results. 
is calculated to be 
, corresponding to 
. The source size can be estimated since both
optically thick and thin lines are available, resulting in >7
,
so that the SiO emission is relatively widespread.
Although the high densities derived above are most likely evidence for
the shock scenario, there remains some uncertainty. MacKay (1995, 1996)
provides an alternative explanation in which silane (
) is
evaporated from icy grain mantles. The silane rapidly forms SiO in the
subsequent chemistry. On basis of the survey data it is difficult to
judge which scenario applies. Higher spatial resolution data are
clearly needed.
The observed ions toward IRS5 have been discussed in detail by de Boisanger et al. (1996) and are only briefly summarized here.
The 3-2 and 4-3 lines of 
and its isotopomers lie at
wavelengths with good atmospheric transmission, and are readily detected
toward the three sources. The main lines are optically thick and the
profile is self-reversed toward W 3(
). The ratio of the two lines
of the more optically thin 
variety is a good density indicator.
For IRS4 it is found that the inferred density is rather high, 
at 55 K. This high value can be due to optical depth
effects (even for 
), to a small source
size, or to calibration errors. The influence on the column density is
not very large, however, and 
is obtained. This value is consistent with that
derived from the 
line. An estimate for the source
size of 17
is found from the optically thick 
line.
De Boisanger et al. (1996) found that the density indicated by
the 4-3/3-2 line ratio toward IRS5 is somewhat higher than that found
from 
but still consistent with the latter result within the
error bars. The column density has been determined using the parameters
of Paper I. No 
has been detected. The source size is
estimated to be 18
.
Toward W 3(
) the 4-3/3-2 ratio indicates a density of 
at 220 K, somewhat lower than found from 
. From
the 
line 
is inferred. This is consistent with
, so that 
. Note that this result does not depend
sensitively on temperature. Decreasing the temperature to 100 K
increases the column density by only 15%. The source size was
calculated from the optically thick 4-3 line. The peak temperature is
found to be more than 15 K, from fits obtained by putting a mask around
the self-absorption. This results in a source size of 
.
The interferometer map of Wink et al. (1994) shows that some
fraction of the emission indeed comes from a compact clump of size 
, but that most of the emission originates in the more extended core
surrounding W 3(
) and W 3(OH).
Much less information is available on 
. There has been an
intensive search for this ion toward IRS5 (see de Boisanger et al.
1996), and the 8-7 line has been detected in the survey toward
W 3(
).
Assuming that the 
stems from the same region as 
toward
IRS4, the upper limit on the 8-7 line implies 
. If the high density found for 
is taken, the
column density decreases to 
.
Toward IRS5, the lines are very weak, but certainly present. It is
interesting that their widths are very small, just as found for the CS
isotopomers. The 8-7/6-5 ratio implies an (uncertain) density of 
, comparable to that found from the 
. The column
density is 
.
The 
line was detected toward W 3(
), but is
unfortunately blended with a methyl formate line. If the methyl formate
excitation is in LTE at 
, the line is not expected to
be very strong, but subthermal excitation may enhance its contribution.
If we assume that 
is the dominant component and use the physical
parameters from Paper I, 
is
obtained.
and 
is interesting because its abundance provides an indirect
determination of that of 
, which cannot be observed at
(sub-)millimeter wavelengths. 
has been detected toward the
Galactic Center in its 
line (Minh et al. 1991b). In
this work, the 
line was searched toward IRS5,
which lies at an energy 
. The non-detection implies
an upper limit on the column density of 
, using a
line strength 
and assuming an excitation
temperature of 50 K.
In the search for ions (de Boisanger et al. 1996), a deep
integration on the 
line at 222 GHz was made toward IRS5,
but the line was not detected, giving an upper limit on the column
density in a 20'' beam of 
.
has been detected toward IRS4 and IRS5 by Turner (1994) in 4
lines, who derived a column density of 
and 
for IRS4 and IRS5 respectively, assuming a source
size of 30''. In this survey the 347740.0 MHz line is detected toward
IRS5, and absent toward W 3(
). The 348115.2 MHz line is
unfortunately blended with 
, and is thus of limited use.
Therefore we adopt the column densities of Turner (1994).
has been searched for toward all three sources by Phillips et
al. (1992) through observations of three lines The detection toward
IRS5 is confirmed in this survey by the 364 GHz line. The total
column density is calculated assuming a high temperature
ortho/para ratio of unity. Using the standard physical parameters,
toward IRS5, neglecting
radiative excitation. However, the 396/364 line ratios observed by
Phillips et al. (1992) indicate that higher densities (
) are more appropriate. This results in
, close to the value
calculated by de Boisanger et al. (1996). The fact that the 364 GHz
line is a factor of 2 stronger in the 15
JCMT beam than in the
20
CSO beam implies a small source size of a few 
.
The non-detection towards IRS4 implies 
) less than
.
In the case of W 3(
) a marginal feature at 364 GHz is present,
with a width much smaller than expected from other lines in this survey
and the data for the 396 GHz line of Phillips et al. (1992). Using the
main beam temperature and the data of Phillips et al. (1992) the
uncertain 
column density is 
.
De Boisanger et al. (1996) have determined the electron fraction
toward IRS5 to be 
, using many
lines and specialized models. For the other two W 3 sources such
detailed information is not available, prohibiting a thorough study of
the electron abundance toward IRS4 and W 3(
). Nevertheless, a
lower limit is easily estimated by adding the abundances of the observed
molecular ions, which is primarily 
. For IRS4 this results in an
abundance 
and for W 3(
) 
.
Only few deuterated molecules have been detected toward the three
sources. One of the most prominent is DCN, which has been seen toward
W 3(
). Unfortunately the 5-4 line is blended with 
,
but it is estimated that 
contributes less than 30% of the
integrated line flux. Assuming standard physical parameters, 
is obtained, giving 
with a factor of four uncertainty. Toward the other two
sources, DCN has not been detected. The amount of deuteration,
[DCN]/[HCN] 
and 
, appears less than
for W 3(
), but can still be the same within the uncertainties.
From the upper limits on the DNC 3-2 line the following upper limits on
N(DNC) are derived: 
and 
toward IRS4 and IRS5 respectively. Toward W 3(
) the DNC 3-2 line
has been marginally detected and using the standard physical conditions
a column density N(DNC) of 
(in a 20
\
beam) is inferred. The [DNC]/HNC] is 
, 
, and 
respectively, similar to the [DCN]/[HCN].
is commonly found in star-forming regions (Wootten et al.
1982). Toward IRS4 the 
line is blended with 
and thus no column density could be derived. Toward IRS5, the 5-4 line
was observed sepearately by de Boisanger et al. (1996). Standard
physical parameters give 
. In
the same way, the beam-averaged column density toward W 3(
) was
found to be 
. This implies 
ratios of 
and 
respectively, again close
to the deuteration values for HCN and HNC.
The high beam-averaged 
column densities found toward the three
sources make an inspection of the HDS upper limits of some interest. If
the same approach is used as for 
, i.e., if 
is assumed, the following HDS upper limits are found from the
line for IRS4, IRS5 and W 3(
) respectively:
; 
; and 
and
; < 0.14 and 
respectively. Unfortunately, these values provide no useful limits on
the deuteration of 
. If the lower excitation temperature of
K is adopted (see Sect. 5.5.5), the limits for IRS5
and W 3(
) both become 
but the ratio
still does not give useful information.
Another deuterated ion is 
for which only upper limits on the
3-2 line were obtained, although a hint of the line is present toward
IRS5 (de Boisanger et al. 1996). Taking the standard physical
parameters, the limits on the column densities are 
;
; and 
for IRS4, IRS5 and
W 3(
) respectively.
Several lines of HDO have been detected toward both W 3(
) and
W 3(OH). The excitation and abundance of this species are described in
a separate paper (Helmich et al. 1996). The inferred 
ratio of 
is lower than that of other species.
Using 
the upper limits on the HDCO column
density for IRS4, IRS5 and W 3(
) are 
and 
respectively. The 
] ratios
are 
; 
and 
respectively. The same method yields 
and
for the 
column densities of IRS4, IRS5
and W 3(
) respectively. The 
ratios are 
; < 0.25 and 
.
H
W 3 IRS4 W 3 IRS5
W 3( 
)
X 
n(H 
) N(X)
n(H 
) N(X) 
n(H 
) N(X)
(K) ( 
cm
) (cm 
) (K) ( 
cm
)
(cm 
) (K) ( 
cm
)
(cm 
)
- - 7 . 0(22) - - 1 . 3(23) - - 9 . 6(22)
C 
O 55 1.0 7 . 3(15) 100 1.0 1 . 3(16) 100 2.0 1 . 0(16)
C 
O 55 1.0 2 . 3(16) 100 1.0 3 . 7(16) 100 2.0 5 . 3(16)
CO - - - - - - 100 2.0 1 . 1(17)
CO - - 1 . 9(19) 
- - 3 . 3(19) 
- - 2 . 6(19) 
CN 55 1.0 3 . 3(14) 100 1.0 1 . 8(14) 220 2.0 1 . 1(14)
HC 
N 55 1.0 3 . 1(12) 100 1.0 2 . 1(12) 220 2.0 6 . 6(12)
H 
CN 55 1.0 1 . 1(13) 100 1.0 5 . 9(12) 220 2.0 1 . 1(13)
HCN - - 8 . 4(14) 
- - 5 . 3(14) 
- -
1 . 2(15) 
HN 
C 55 1.0 3 . 3(12) 100 1.0 3 . 2(12) 220 2.0 1 . 5(12)
HNC - - 2 . 0(14) 
- - 1 . 9(14) 
- -
9 . 0(13) 
HC 
N 55 1.0 <2 . 5(13) 100 1.0 <9 . 8(12) 220 2.0 1 . 5(13)
HNCO 55 
- <2 . 8(13) 100 
- 4 . 4(13) 53 
- 4 . 8(14)
CCH 55 1.0 1 . 9(14) 100 1.0 1 . 2(14) 220 2.0 1 . 3(14)
CH 
C
H 25 
- 7 . 7(14) 50 
- <1 . 0(14) 63 
- 6 . 1(14)
H 
CO - - - - - - 220 2.0 5 . 1(12)
55 1.0 8 . 0(13) 100 1.0 8 . 0(13) 220 2.0 4 . 0(14)
28 
- 4 . 2(14) 47 
- 1 . 6(14) 265 
- 8 . 8(15)
CH 
CN 55 1.0 <1 . 2(13) 100 1.0 <9 . 4(12) 120 4.0 2 . 7(13)
- - - - - - 193 
- 2 . 0(15)
- - - - - - 141 
- 6 . 7(14)
C 
S 55 1.0 6 . 4(12) 100 1.0 1 . 4(12) 220 2.0 1 . 5(13)
CS 55 1.0 6 . 6(12) 100 1.0 2 . 5(12) 220 2.0 2 . 0(13)
C 
S 55 1.0 1 . 9(13) 100 1.0 1 . 0(13) 220 2.0 4 . 3(13)
CS - - 4 . 0(14) 
- - 2 . 0(14) 
- -
1 . 0(15) 
S 
O - - - 200 10 1 . 0(13) - - -
SO - - - 200 10 5 . 3(13) 220 2.0 4 . 3(13)
SO - - - 200 10 8 . 0(13) 220 2.0 5 . 6(13)
SO 80 2.0 1 . 4(14) - - 5 . 0(15) 
- -
1 . 3(15) 
102 
- <4 . 4(13) 147 
- 3 . 0(14) 179 
- 2 . 4(14)
102 
- 1 . 5(14) 154 
- 5 . 3(15) 184 
- 1 . 0(15)
80 2.0 4 . 4(13) 100 1.0 1 . 3(13) 100 2.0 1 . 5(14)
55 
- 1 . 1(14) 100 
- 2 . 0(14) 220 
-
1 . 0(15)
OCS 55 1.0 <5 . 0(13) 150 2.0 8 . 9(13) 220 2.0 2 . 3(14)
SiO - - - - - - 220 2.0 3 . 0(12)
SiO 55 1.0 3 . 7(12) 250 10 5 . 3(12) - -
5 . 9(13) 
HC 
O
55 10 6 . 7(11) - - - 220 2.0 4 . 3(11)
H 
CO
55 10 6 . 3(12) 100 1.0 3 . 3(12) 220 2.0 3 . 9(12)
HCO 
- - 3 . 8(14) 
- - 1 . 9(14) 
- -
2 . 3(14) 
HCS 
55 1.0 <2 . 0(13) 100 1.0 6 . 3(11) 220 2.0
1. 7(13)
HOCO 
- - - 50 
- <2 . 0(12) - - -
HCNH 
- - - 100 1.0 <1 . 5(13) - - -
H 
O
55 1.0 <4 . 6(13) 100 5.0 1 . 2(14) 220 2.0 1 . 0(14)
DCN 55 1.0 <3 . 6(12) 100 1.0 <2 . 6(12) 220 2.0 9 . 0(12)
DNC 55 1.0 <1 . 4(12) 100 1.0 <7 . 7(11) 220 2.0 4 . 3(11)
DCO 
- - - 100 1.0 4 . 8(11) 220 2.0 9 . 5(11)
HDS 55 
- <1 . 0(13) 100 
- <2 . 7(13) 220 
-
<8 . 1(13)
N 
D
55 1.0 <3 . 1(11) 100 1.0 <3 . 3(11) 220 2.0 <2 . 7(11)
HDCO 55 
- <3(12) 100 
- <4(12) 220 
- <1(13)
CH 
OD 55 
- <3(13) 100 
- <4(13) 220 
- <1(14)
. . .
; see text for derived physical parameters
from observed line ratios of each molecule.
Inferred rotational temperature rather than kinetic temperature.
The rotation diagram method has been used to determine the listed column
densities for these species.
Value derived from 
lines, see text.
Column density of main isotopomer determined from optically thin
isotopomers.
Apart from the torsional modes of 
, several other lines
originating in vibrationally-excited levels were found in the W 3
survey. Two J= 4-3 lines within the first vibrational level of the
bending mode 
of HCN were detected toward W 3(
).
This mode lies at much higher energies, approximately 1000 K, than the
torsional modes of 
, so that the levels are unaccessible in most
molecular cloud environments. Vibrationally excited HCN was first
detected by Ziurys & Turner (1986) in Orion. In principle there are
two ways to populate such high states. The first interpretation,
favoured by Ziurys & Turner, is through absorption of 
photons
into the 
bending mode. In the case of W 3(
),
these photons can be provided by the embedded late O/early B star which
heats the dust in the inner part of the envelope and the compact clump
surrounding the star. Depending on the actual physical conditions and
geometry, the rotational level populations in the ground vibrational
state can also be controlled by this infrared pumping (see e.g.,
Hauschildt et al. 1993). Note that the continuum from W 3(OH) is not
able to pump the HCN around W 3(
), as Turner & Welch
(1984) already concluded.
The second mechanism for populating the vibrationally excited levels is
through collisions at sufficiently high densities and temperatures.
Wink et al. (1994) estimate a density of 
on a 1
scale, whereas Wilner et al. (1995) obtain 
within 0.5''. Although the optically thick dust emission
prevents any submillimeter lines from <1'' to be observed, the
density in a 1-2'' region is obviously very high and close to the
critical density for excitation. More detailed modeling is needed to
assess which process dominates in the case of W 3(
).
The J = 7-6 line within the 
level of CS was detected in
three independent spectra toward W 3(
). This vibrational level
lies at approximately 1900 K (
) and is thus even more
difficult to populate than that of HCN. This suggests that close to
W 3(
), either the dust must be very hot or the conditions quite
extreme (see also Helmich et al. 1996).
CCH has its lowest vibrational level at 
or 
(Hsu et al. 1993; Kanamori & Hirota 1988). At these
wavelengths the infrared radiation is intense enough to pump this
vibrational level. Black (private communication) predicts the lines
within the 345 GHz window to lie at 346248, 346929, 348974 and
349650 MHz. However, no CCH lines within this vibrational level are
found in the survey. This suggests that the CCH abundance in the inner
region is lower than those of HCN and CS. CCH is also unique in that it
has its first electronic state at low energies (
).
As in other surveys, some lines remain unidentified, but the fraction is small. The U-lines are generally quite weak, but are thought to be true detections. Because the data are taken in double side-band mode, there are two possible frequencies listed, unless the U-line turns up in more than one spectrum with different local oscillator settings. In those cases, a definite determination of the frequency can be made and only one value is given. Most of the U-lines (7) are found toward IRS5.
|
|  |  | W | ||
| (MHz) | (K) | (km s ) | (K km s ) | ||
|
W 3 IRS4 | |||||
| 342794 | 0.37 | 5 | . | 96 | 2.35 |
| 342803 | 0.58 | 3 | . | 97 | 2.43 |
| W 3 IRS5 | |||||
| 338855/342255 | 0.23 | 5 | . | 21 | 1.25 |
| 338966/342106 | 0.083 | 9 | . | 47 | 0.84 |
| 339098/342012 | 0.11 | 10 | . | 7 | 1.24 |
| 339209/341901 | 0.079 | 4 | . | 01 | 0.34 |
| 339220/341890 | 0.073 | 4 | . | 19 | 0.33 |
| 342794 | 0.48 | 4 | . | 12 | 2.09 |
| 342803 | 0.49 | 4 | . | 17 | 2.19 |
W 3( ) | |||||
| 362098 | 0.24 | 6 | . | 53 | 1.67 |
| 231310/233990 | 0.093 | 7 | . | 27 | 0.72 |
| 336497/339897 | 0.22 | 2 | . | 76 | 0.65 |
The fact that the three sources have different chemical characteristics
provides hints on the identity of the carriers of the lines. Small,
non-saturated molecules are mostly found toward IRS4, so there is a
large probability of identifying the two U-lines toward this source with
a simple molecule. Since the same U-lines are also found toward IRS5,
the molecule may involve sulfur. Indeed, the other U-lines toward IRS5,
which are not found toward the other sources, most likely contain sulfur
or silicon. The U-lines toward W 3(
) can be due to any species,
but most likely originate from some complex organic molecule.
The U-lines were compared with those found in surveys for other sources but there are no U-lines in common with either the Orion surveys of Sutton et al. (1995) and Schilke et al. (1996), the Sgr-B2 survey of Sutton et al. (1991) or the IRAS 16293-2422 survey of Blake et al. (1994).
fit parameters and beam-averaged column
densities