In the previous section, a detailed description of the observations and analysis for each individual species has been given. We discuss here the results in general, both for the physical structure and the chemical state of the three sources.
In Paper I, the physical parameters were derived primarily from the
observed excitation under the assumption that a single physical
component is present within the 15'' beam (0.16 pc). We have analyzed
here the excitation of many more species and can compare the results
with those found from
. At the same time, it provides
information on the question of multiple components within the beam. In
this respect, it is useful to compare the observed scale in W 3 with
that of structures found in the more nearby Orion giant molecular cloud,
which is at 450 rather than 2300 pc. From observations of the extended
Orion ridge component, it appears that variations in physical and
chemical parameters occur on scales of a few arcmin, corresponding to
(Ungerechts et al. 1995;
Schilke et al. 1992). Thus, for the surrounding, more quiescent molecular
cloud material, the linear scales probed in W 3 are comparable and the
assumption of a single component may be reasonable. However, in the
Orion/KL region around the massive young stellar object IRc2,
significant changes occur on scales of < 10'', corresponding to
. Structures of similar linear size in the W 3 region will be
strongly diluted in the 15'' beam.
For IRS4, the deviations from the picture of a single extended,
quiescent warm cloud are the smallest. The cloud seems to be well
described by a density and
, as derived from
. The
excitation of other species such as CS, CN, CCH and HNC results in
similar parameters. However, there is some evidence that warmer
material is present in the beam. Weak evidence comes from the fact that
the CO main beam brightness temperatures and the
line ratios
indicate temperatures of order
. The strongest case is formed
by the
rotational temperature
.
Sulfur-bearing molecules have also been found in other regions to be
associated with elevated temperatures and shocks (Ziurys 1991).
There are several ways to create elevated temperatures in the IRS4
region. The most likely are through shocks or through heating by
ultraviolet radiation. In star-forming regions shocks are often
associated with outflowing material, but IRS4 does not show any clear
sign of an outflow in its line profiles (see Sect. 5.1). Recently,
``hyper-compact'' objects (e.g. Claussen et al. 1994) have
been discovered in giant molecular clouds. These are small (<1'')
clumps of ionized material seen in VLA measurements. The ionization
presumably originates from embedded B-stars, but is kept to small scales
due to very high densities close to the star. Therefore it will only
have an impact on very small scales within the molecular cloud. In
Paper I, we postulated the picture that IRS4 is located on the far side
of the molecular cloud and has already broken free on one side,
illuminating the cloud from behind. In this scenario, the ultraviolet
radiation from the newly formed star can heat the gas and set up a
photon-dominated region (PDR) on the back side. In addition, it will
drive an ionization front into the dense material. At the interface
between the expanding ionization front and the molecular cloud, a shock
will be present, which can also heat the gas and return grain mantle
material. The shocked gas can be quite extended, since the H II
region itself already has a size of several arcsec (Colley
1980). The PDR will extend over similar scales or larger.
This interpretation can be distinguished from the hyper-compacts
scenario through interferometer observations of species like
.
In the case of IRS5, more than one physical component is certainly
present in the beam, since it is well known that the source has a strong
outflow (see HMMT 1994 for the most recent results). In addition, the
extremely deep self-absorption in the CO profiles shows that there must
be a temperature and density gradient or stratification. The depth of
the absorption in the lower-J lines indicates that low temperature
() and/or low density gas is present on the outside. The fact
that the 6-5 line still shows self-absorption (HMMT 1994), but not the
9-8 line (Boreiko & Betz 1991), suggests that a range of conditions is
present. The latter authors give an upper limit on the temperature of
the absorbing gas of 80 K, whereas HMMT find an excitation temperature
of 26 K. From the smaller self-absorption in the
line
and its absence in the
lines, the column density of the
absorbing gas is limited to
. Since this
is a factor of 6 smaller than the column density of the emitting gas and
since our submillimeter observations are biased toward the warmer and
denser gas, our results are not expected to refer to the cold component.
Although gradients are likely to be present, the excitation analysis
shows that most molecules (including CS, , HCN) are located in
fairly extended gas with
and
. Part of this agreement stems from the fact
that most lines in the 230 and 345 GHz windows have critical densities
of order a few
, so that it is no surprise that
most detected species give a similar result. Observations of higher
excitation lines are needed to determine whether the density is fairly
homogeneous, or whether a gradient or power law distribution is present,
as would be expected for a collapsing cloud core.
A physically distinct region appears to be probed by the sulfur-bearing
molecules, at least by those containing oxygen. In Paper I, it was
already found that the emission comes from a warmer region, which
was tentatively connected with the outflow. The rotational temperature
has been better constrained in this work, but the same conclusion
remains (see Sect. 5.5.3). In addition, it was found from the
strengths of the optically thick lines of SO and
that the source
size is
. Just as in IRS4, the chemical composition of this
material seems disjunct from that of the bulk of the gas. In Orion
IRc2,
turned out to be a major coolant of the gas in this
frequency range (Groesbeck 1994). This is not the case for IRS5. Not
only is the column density much lower than that toward IRc2, but the
line width is smaller as well. A crude estimate shows that the total
emission in the 345 GHz range is less than half of that of CO,
putting IRS5 somewhere in between IRc2 and the Orion-S source in this
respect (Groesbeck 1994).
may be formed from the large
amount of
in the same gas.
The region probed by SiO appears to be even denser with
. Similarly high densities and even
higher temperatures have been found for SiO toward Orion IRc2, which has
brightness temperatures in interferometer observations of
(Wright et al. 1995). This suggests disruption of grains by strong
shocks and rapid formation of SiO in the warm post-shock gas. Note,
however, that in contrast with Orion/KL, the SiO and
lines are
not noticeably broader toward IRS5 than those of other species. Also,
the SiO is not kinematically displaced in contrast with L 1448
(Guilloteau et al. 1992) and NGC 1333 IRAS4 (Blake et al. 1995).
Interferometric observations of both SiO and
are needed to
properly investigate these hot and dense regions.
The fact that no rotational lines within the first vibrational levels of HCN, CS and/or CCH have been detected toward IRS5 implies that the hot and dense gas does not contain large quantities of these molecules.
At least two different components are observed toward W 3() within
the 15'' JCMT beam. For many species, a significant fraction of the
emission comes from the warm
and dense (
) core (>20'') in which both W 3(
) and W 3(OH)
are located.
The second component is the more compact ``hot core'' type region, which
is warmer () and denser (
) than the surrounding envelope. The saturated organic
molecules such as
(Wink et al. 1994), part of the
and
, and probably most of the dimethyl ether and methyl formate are
found here as well, based on their rotational temperatures. The
unresolved continuum source of <1 '' (Wilner et al. 1995) is too
diluted in our beam to result in detectable emission. In addition, at
submillimeter wavelengths the continuum is optically thick, so that the
line emission must originate from a larger region. This strong
continuum radiation can, however, affect the excitation of molecules
such as HNCO and HDO (Helmich et al. 1996), and lead to large
populations in vibrationally excited levels of molecules such as HCN and
CS.
The , SO and SiO lines again indicate high temperature and
densities. Most likely, these species are located in the outflow traced
by the
masers (Alcolea et al. 1992). This region
probably contains the
found by Phillips et al. (1992).
Surveys like the one presented here are ideal tools to establish the
beam-averaged composition of the gas, since the excitation effects can
largely be taken out. As discussed in Sect. 4.3, it is more difficult
to make the final step toward abundances, which can be directly compared
with models. The beam-averaged column densities are obtained from
, but this refers to the cloud as a whole whereas some
molecules may be more localized. The excitation analysis only gives an
estimate of the local
number density, not of its column density.
For some species such as
in IRS5 or
in W 3(
), the
local abundances may be more than two orders of magnitude higher than
the beam averaged values listed in Table 5 (click here) (see Sect. 5.4.3). These
abundances should therefore be interpreted with care.
| W 3 | W 3 | W 3 | Extended | Compact | ||||||||||
IRS4 | IRS5 | (![]() | Ridge![]() | Ridge![]() | |||||||||||
H | 0 | . | 5 | 0 | . | 5 | 0 | . | 5 | 0 | . | 5 | 0 | . | 5 |
CO | 1 | . | 35(-4) | 1 | . | 35(-4) | 1 | . | 35(-4) | 1 | . | 35(-4) | 1 | . | 35(-4) |
CN | 2 | . | 4(-9) | 6 | . | 9(-10) | 5 | . | 7(-10) | 1 | . | 0(-8) | 1 | . | 9(-9) |
HCN | 6 | . | 0(-9) | 2 | . | 0(-9) | 6 | . | 3(-9) | ||||||
HNC | 1 | . | 4(-9) | 7 | . | 3(-10) | 4 | . | 7(-10) | ||||||
HC![]() | <1 | . | 3(-10) | <3 | . | 8(-11) | 7 | . | 8(-11) | 1 | . | 5(-8) | |||
HNCO | <2 | . | 0(-10) | 1 | . | 7(-10) | 2 | . | 5(-9) | ||||||
CCH | 1 | . | 4(-9) | 4 | . | 6(-10) | 6 | . | 8(-10) | ||||||
CH![]() ![]() | 5 | . | 5(-9) | <3 | . | 9(-10) | 3 | . | 2(-9) | 2 | . | 3(-8) | 3 | . | 1(-9) |
![]() | 5 | . | 7(-10) | 3 | . | 1(-10) | 2 | . | 1(-9) | <2 | . | 6(-8) | 9 | . | 1(-8) |
![]() | 3 | . | 0(-9) | 6 | . | 2(-10) | 4 | . | 6(-8) | 1 | . | 9(-8) | 9 | . | 4(-7) |
CH![]() | <1 | . | 4(-10) | <7 | . | 2(-11) | 1 | . | 4(-10) | 1 | . | 2(-8) | |||
![]() | - | - | 1 | . | 0(-8) | <5 | . | 6(-9) | 4 | . | 7(-8) | ||||
![]() | - | - | 3 | . | 5(-9) | <1 | . | 1(-8) | 7 | . | 8(-8) | ||||
CS | 2 | . | 9(-9) | 7 | . | 7(-10) | 5 | . | 2(-9) | 3 | . | 3(-8) | 2 | . | 5(-8) |
SO | 1 | . | 0(-9) | 1 | . | 9(-8) | 6 | . | 8(-9) | 1 | . | 1(-8) | 6 | . | 3(-7) |
![]() | 1 | . | 1(-9) | 2 | . | 0(-8) | 5 | . | 2(-9) | 4 | . | 6(-9) | 4 | . | 1(-7) |
![]() | 3 | . | 1(-10) | 5 | . | 0(-11) | 7 | . | 8(-10) | 3 | . | 4(-9) | 3 | . | 1(-9) |
![]() | 7 | . | 9(-10) | 7 | . | 7(-10) | 5 | . | 2(-9) | ||||||
OCS | 3 | . | 6(-10) | 3 | . | 4(-10) | 1 | . | 2(-9) | <1 | . | 0(-8) | 6 | . | 9(-9) |
SiO | 2 | . | 6(-11) | 2 | . | 0(-11) | 3 | . | 1(-10) | <1 | . | 5(-9) | 1 | . | 9(-8) |
HCO![]() | 2 | . | 7(-9) | 7 | . | 3(-10) | 1 | . | 2(-9) | 1 | . | 1(-8) | 2 | . | 4(-9) |
HCS![]() | <1 | . | 4(-10) | 2 | . | 4(-12) | ![]() | . | 9(-11) | 5 | . | 9(-10) | 1 | . | 0(-10) |
HOCO![]() | - | <7 | . | 7(-12) | - | ||||||||||
HNCH![]() | - | <5 | . | 8(-11) | - | ||||||||||
H![]() ![]() | <3 | . | 3(-10) | 4 | . | 6(-10) | 5 | . | 2(-10) | ||||||
| . | . | . | . | . |
The resulting abundances are listed in Table 5 (click here). They reinforce the main
conclusion from Paper I that there are large chemical differences
between the three sources which cannot be ascribed to excitation
effects. Toward IRS4, mostly simple molecules are found. The
abundances of these species are low toward IRS5, except those of sulfur-
and silicon-containing species. Toward W 3(), the abundances of
complex organic species such as
,
, dimethyl ether and
methyl formate are at least 1-2 orders of magnitude larger than those
found toward the other two sources. For comparison, the abundances
found by Sutton et al. (1995) for the Orion extended ridge and the
compact ridge scaled to the same CO abundances are included. In the
following, some aspects of the chemistry in the three sources will be
discussed in more detail.
Toward IRS4 the abundances of most molecules listed in Table 5 (click here) are
expected to be reliable, since their emission fills the 230 and 345 GHz
beams to first order. Moreover, there seems to be no indication of a
density gradient for these species. Only the SO, and perhaps SiO
pose some problems, if they arise in a more localized, hotter (shocked)
region. Their listed abundances should be regarded as lower limits.
The abundances of unsaturated radicals and molecules like CN, and
are large toward IRS4, suggesting that these molecules are
present mostly in the quiescent gas and avoid the shock or ``hot core''
type conditions. This is consistent with observations of the low mass
young stellar object IRAS 16293-2422, where CN and
were also
found to trace the colder, more quiescent outer part of the circumbinary
envelope (van Dishoeck et al. 1995). CN and
are known to be
prominent in the outer zones of photon-dominated regions (PDR) such as
the Orion Bar (Jansen et al. 1995;
Sternberg & Dalgarno 1995;
Simon et
al. 1995; Fuente et al. 1995). Indeed, the observed
abundances of simple molecules in IRS4 are very similar to those found
in the Orion Bar (see Helmich et al. 1997). The extended ridge in
Orion also shows enhancements in these species. The ultraviolet
radiation from IRS4 (and any other stars in the W 3 region) could create
a PDR at the back of the cloud, resulting in large abundances of these
species.
The abundances toward IRS5 are more difficult to interpret since there
certainly is some stratification of material in the beam. Nevertheless,
there is no doubt that the (beam-averaged) abundances of and SiO
in this source are larger by at least an order of magnitude compared
with IRS4. Surprisingly, however, the abundances of sulfur-bearing
molecules containing carbon (CS,
) appear lower than those found
toward IRS4. Their line widths are also narrower, suggesting that these
species are located in the outer, more quiescent part of the envelope.
The proposed connection of
and SiO with the outflow in IRS5 needs
to be tested further through interferometer observations of these
species. If the column density derived by Mitchell et al. (1991) for
the outflow is used, the abundances are increased by up to two orders of
magnitude.
Although IRS5 has a powerful outflow and emits copious infrared photons,
apparently only a small fraction of the envelope has been affected,
since the abundances of many species, including HCN and , are
lower compared with IRS4 and W 3(
). In Paper I, we suggested that
at least some of these molecules are still present in the ice mantles on
the grains. This case is particularly strong for
, where the
amount of solid
as measured by Allamandola et al. (1992) is 4
orders of magnitude larger than that of gas phase
. The amount
of depletion for many other molecules is still unknown, however,
although most of the CO is certainly located in the gas phase (Tielens
et al. 1991). Observations with the Infrared Space Observatory (ISO)
will provide a complete inventory of the ice-mantle composition, which
together with this work can be used to determine the relative amounts
present in the gas and on the grains (cf. van Dishoeck et al. 1996).
Toward W 3() the largest complication is formed by the fact that
the relative abundances of the molecules in the more extended core and
in the compact clump are poorly determined. Nevertheless, the high
temperature, the high methanol abundance, and the large column densities
of methyl formate, dimethyl ether and OCS are clear signs of a ``hot
core'' chemistry within the beam, in which molecules such as
are released from the grains and drive a rapid gas-phase chemistry
leading to complex organic molecules. This chemistry has been modelled
e.g. by Charnley et al. (1992), and from their work it is obvious that
at least
must have passed since the onset of methanol release.
The amount of methanol relative to methyl formate and dimethyl ether is
small, an order of magnitude smaller than found in e.g. the Orion
compact ridge (see Table 5 (click here)), and therefore difficult to account for in
chemical models. The situation is worse in G34.3 (Millar private
communication) where the methanol column density equals the dimethyl
ether and methyl formate column densities. High optical depths can
influence the methanol abundance severely, but as shown above it is not
very likely to be significant here. More detailed chemical modeling
will be presented in a later paper (Helmich et al. 1997).
The HCN over HNC ratio has been used to determine time scales for the chemical evolution (Goldsmith et al. 1986; Irvine et al. 1987; Schilke et al. 1992). It is beyond the scope of this work to perform a similar analysis here for the W 3 sources, but it is worthwhile to note that the values for the three W 3 sources lie somewhere between those of Orion IRc2 and the Ridge.
Only few deuterated species have been found in this work. The simplest
molecules like , DCN and DNC have been detected and these show
indeed signs of some deuterium fractionation suggesting that the cloud
had lower temperatures in its past. On the other hand, the observed
amount of deuteration is less than the values of 0.01-0.1 found in cold
clouds such as TMC-1, and more comparable to that found in other warm
clouds (see e.g. Table 2 (click here) of van Dishoeck et al. 1993a for an
overview). It is also much lower than found in the low-mass
star-forming region IRAS 16293-2422 (van Dishoeck et al. 1995).
Together with the relatively low deuteration for HDO (Helmich et al.
1996) this can be interpreted as an indication that the cloud never went
through a very cold (
) phase. The limits on most other species
such as HDCO and
are not very stringent.
Ratio | W 3 IRS4 | W 3 IRS5 | W 3(![]() | ||||||
DCN/HCN | < 4 | . | ![]() | < 5 | . | ![]() | 7 | . | ![]() |
DNC/HNC | < 7 | . | ![]() | < 4 | . | ![]() | 4 | . | ![]() |
DCO![]() ![]() | - | 2 | . | ![]() | 4 | . | ![]() | ||
HDS/![]() | < 9 | . | ![]() | < 0 | . | 14 | < 8 | . | ![]() |
HDCO/![]() | < 3 | . | ![]() | < 5 | . | ![]() | < 2 | . | ![]() |
CH![]() ![]() | < 7 | . | ![]() | < 0 | . | 25 | < 1 | . | ![]() |
HDO/![]() | - | - | ![]() | ||||||
| . | . | . |