A&A 381, 1026-1038 (2002)
H. Roberts 1, - G. A. Fuller 1 - T. J. Millar 1 - J. Hatchell 2 - J. V. Buckle 1,
1 - Department of Physics, UMIST, PO Box 88, Manchester M60 1QD, UK
2 - Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany
Received 30 July 2001 / Accepted 5 October 2001
We present observations of [HDCO]/[H2CO] and [DCN]/[HCN] ratios towards a selection of low-mass protostellar cores in three different star forming regions. The best fit to the observed [HDCO]/[H2CO] ratios is 0.05-0.07, similar to the values observed towards the dark clouds, TMC-1 and L134N. [DCN]/[HCN] ratios are 0.04, higher than those seen in TMC-1, around the low-mass protostar IRAS16293 and the Orion Hot Core, but similar values to the Orion Compact Ridge and L134N. We compare our results with predictions from detailed, chemical models, and to other observations made in these sources. We find best agreement between models and observations by assuming that interaction between gas phase molecules and dust grains has impacted on the chemistry during the cold pre-collapse phase of the cloud's history. The abundance of deuterated species indicates that the dense gas out of which a low mass protostar forms, evolves and collapses on a timescale of 50000 years. We find no marked difference between molecular D/H ratios towards different regions, or between Class 0 and Class I protostars. However, the striking difference between the [DCN]/[HCN] ratios we have measured and those previously observed towards Hot Molecular Cores leads us to suggest that there are significant evolutionary differences between high and low mass star forming regions.
Key words: ISM: molecules - ISM: clouds - ISM: abundances - stars: formation
Observations of deuterated molecules have become an important tool in the study of interstellar chemistry. Although the underlying (or cosmic) D/H ratio is low (10-5), the formation of deuterated molecules is preferred at low temperatures (80 K) and leads to a high degree of fractionation in cold, dark clouds. For example, in the quiescent dark cloud, TMC-1, molecular D/H ratios, including [HDCO]/[H2CO] and [DCN]/[HCN] are observed to be >10-2.
Enhanced molecular D/H ratios are also observed in hot molecular cores (HMC's), clumps of hot, dense gas, usually associated with high mass star formation. The temperatures of these cores (typically 70-150 K) should be high enough to preclude the enhancement of molecular D/H ratios through gas-phase reactions. However, the ratios which have been measured are generally 10-3 (e.g. Hatchell et al. 1998, 1999), lower than TMC-1, yet still enhanced over the cosmic value. It is now generally accepted that these ratios have been preserved from an earlier, colder phase of the cloud's history in the ice-mantles of dust grains. Once some heating event, such as the formation of a star or the passage of a shock, heats the grains sufficiently to evaporate their mantles, the D/H ratios can survive for 104 yrs in the hot gas (Rodgers & Millar 1996).
In cold cores which are forming low mass stars, we might expect a situation intermediate between hot cores and dark clouds. To date, the only survey of deuterated molecules in a low-mass star forming region has been that of IRAS16293-2422 (hereafter, IRAS16293), a class 0, proto-binary system in Oph, by vanDishoeck et al. (1995). However, their survey revealed discrepancies in the levels of fractionation of different molecules, with over 10% deuteration seen in species such as HDCO and HDS, yet only a few percent in species such as DCN. While the [DCN]/[HCN] ratio is similar to that observed in TMC-1, the [HDCO]/[H2CO] ratio is at least twice as high. Neither the very large [HDCO]/[H2CO] or [HDS]/[H2S] ratios can be explained by a standard gas-phase chemistry.
We wished to confirm whether these high D/H ratios were a general feature of low-mass star formation, or particular to IRAS16293. Therefore we have carried out a survey to measure both the [HDCO]/[H2CO] and [DCN]/[HCN] ratios in the dense gas associated with young protostars (`protostellar cores') in three different star forming regions. Our sources are listed in Table 1.
|[ ]||[ ' '']||(kms-1)||(K)|
Sections 2 and 3 describe the observations and data reduction techniques, presenting the resulting column densities and molecular D/H ratios, Sect. 4 summarises and discusses these results. In Sect. 5 we describe the chemical models and compare model predictions with the observations. Section 6 compares these results with those from previous observations of high-mass star forming regions. Throughout, we adopt the conventions; "N(ABC)'' for the column density of molecule ABC, "[ABC]'' for N(ABC)/N(H2), i.e. the fractional abundance of molecule ABC, and "fractionation of XD'' for [XD]/[XH].
The observations were carried out using the NRAO12 m radio telescope, at Kitt Peak, Arizona, in November 1999, Table 2 lists the observed transitions.
|H2CO 21,1-11,0 (ortho)||150.498||15.7||6.468 10-5|
|H2CO 51,4-51,5 (ortho)||72.409||45.8||8.011 10-7|
|H213CO 21,2-11,1(ortho)||137.450||15.1||4.928 10-5|
|HDCO 21,1-11,0||134.285||9.3||4.587 10-5|
|HDCO 40,4-30,3a||256.585||21.5||4.736 10-4|
|HCN 1-0 (triplet)||88.632||2.96||2.426 10-5|
|H13CN 1-0(triplet)||86.340||2.88||2.224 10-5|
|HC15N 1-0||86.055||2.87||2.203 10-5|
|DCN 1-0 (triplet)||72.415||2.42||1.312 10-5|
|DCN 2-1(multiplet)b||144.828||7.25||1.259 10-4|
|c-C3HD 22,0-11,1c||137.455||6.27||5.335 10-5|
|a Observed towards only three sources, due to high
b Observed by Buckle & Fuller (2001).
c Tentative identification of lines seen in the H213CO band.
The 3 mm receivers were used to observe the HCN1-0 and DCN1-0 transitions, the H2CO51,4-51,5 transition frequency being covered by the DCN1-0 band, the 2 mm receivers were used for the H2CO21,1-11,0, HDCO21,1-11,0 and H213CO21,2-11,1 transitions and the 1 mm receivers were used for the HDCO40,4-30,3 line. The MAC spectrometer was in 2 IF mode with 16384 channels. The H13CN and HC15N transitions were observed simultaneously, with the 3 mm receiver, using 4 IF mode with an offset of 285 MHz between the pairs of filter banks, though in the end, the HC15N lines were not above the level of the noise. In all cases the spectra were observed using 24 kHz channels, giving a velocity resolution of 0.04 kms-1, except for the DCN2-1 spectra, observations of which are described in Buckle & Fuller (2001).
The weather was clear, with typical sytem temperatures of 180-300 K for the 3 mm receiver and 200-350 K for the 2 mm receiver, while for the brief time we used the 1 mm receiver, the system temperature rose to 1000 K. Pointing was checked every two hours or so and errors found to be 5-6''.
All data were calibrated by the usual chopper wheel method and corrected for , the efficiency at which the source couples to the main diffraction beam, to give . In all cases we have assumed a filling factor of 1, i.e. we have assumed that the source fills the main diffraction beam of the telescope. This may not be strictly true for our low-frequency transitions (the beamwidth of the 12 m telescope is 90'' at 70 GHz), however, as we are primarily interested in the ratios of our column densities, any errors introduced by this assumption should cancel out.
The spectra we obtained are shown in Figs. 1 and 2.
|Figure 1: Spectra of H2CO (left), H213CO (centre) and HDCO (right) observed towards the sources in our survey. The 31,2-22,1 transition of c-C3HD has been tentatively identified in the H213CO spectra towards B5IRS1, L1448 mms, L1448 NW and L1527.|
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|Figure 2: Spectra of HCN (left), H13CN (centre) and DCN (right) obtained towards the sources in our survey. The approximate expected position of the H2CO 51,4-51,5 transition is indicated on each DCN spectra.|
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|H2CO||B5IRS1||0.66 (0.02)||0.771 (0.077)|
|21,1-11,0||L1448 mms||1.96 (0.03)||2.013 (0.181)|
|L1448 NW||1.53 (0.02)||3.962 (0.198)|
|HH211||1.02 (0.01)||3.152 (0.126)|
|IRAS03282||0.96 (0.03)||0.934 (0.100)|
|L1527||1.02 (0.04)||1.120 (0.123)|
|L1551IRS5||1.08 (0.01)||2.722 (0.136)|
|RNO43||0.88 (0.03)||0.873 (0.096)|
|HH111||1.33 (0.03)||1.342 (0.161)|
|L1448 NW||0.99 (0.09)||0.164 (0.061)|
|HH211||0.48 (0.06)||0.166 (0.071)|
|L1527||0.44 (0.05)||0.100 (0.027)|
|L1551IRS5||0.52 (0.07)||0.080 (0.016)|
|HH111||0.97 (0.13)||0.095 (0.047)|
|HDCO||B5IRS1||0.67 (0.10)||0.054 (0.034)|
|21,1-11,0||L1448 mms||1.28 (0.14)||0.126 (0.077)|
|L1448 NW||0.72 (0.04)||0.405 (0.069)|
|HH211||0.44 (0.03)||0.149 (0.021)|
|IRAS03282||0.52 (0.10)||0.073 (0.030)|
|L1527||0.44 (0.02)||0.237 (0.026)|
|L1551IRS5||0.71 (0.04)||0.210 (0.036)|
|HH111||0.60 (0.08)||0.110 (0.051)|
|c-C3HD||B5IRS1||0.46 (0.06)||0.049 (0.021)|
|22,0-11,1||L1448 mms||1.00 (0.12)||0.100 (0.040)|
|L1448 NW||0.75 (0.10)||0.106 (0.046)|
|L1527||0.35 (0.03)||0.153 (0.021)|
|HCN||B5IRS1||0.43 (0.01)||0.542 (0.705)|
|1-0||L1448 mms||1.48 (0.03)||1.747 (0.140)|
|L1448 NW||1.40 (0.01)||4.576 (0.275)|
|HH211||1.01 (0.01)||2.127 (0.149)|
|IRAS03282||0.69 (0.02)||0.688 (0.083)|
|L1527||0.48 (0.02)||0.391 (0.070)|
|L1551IRS5||0.74 (0.01)||0.877 (0.096)|
|RNO43||0.99 (0.05)||0.584 (0.123)|
|HH111||0.89 (0.05)||0.949 (0.247)|
|1-0||L1448 mms||1.11 (0.07)||0.299 (0.057)|
|L1448 NW||1.27 (0.08)||0.460 (0.069)|
|HH211||0.83 (0.10)||0.206 (0.037)|
|IRAS03282||0.91 (0.07)||0.129 (0.036)|
|L1527||0.41 (0.06)||0.092 (0.022)|
|L1551IRS5||0.62 (0.07)||0.114 (0.024)|
|DCN||B5IRS1||0.49 (0.04)||0.200 (0.030)|
|1-0||L1448 mms||1.27 (0.06)||0.681 (0.075)|
|L1448 NW||1.01 (0.05)||0.843 (0.084)|
|HH211||0.92 (0.07)||0.397 (0.052)|
|IRAS03282||0.71 (0.04)||0.267 (0.032)|
|L1527||0.39 (0.03)||0.180 (0.023)|
|L1551IRS5||0.73 (0.05)||0.285 (0.037)|
|DCN||B5IRS1||0.57 (0.05)||0.131 (0.042)|
|2-1||L1448 mms||1.07 (0.09)||0.476 (0.071)|
|L1448 NW||1.03 (0.09)||0.515 (0.041)|
|HH211||0.81 (0.06)||0.224 (0.029)|
|IRAS03282||0.64 (0.05)||0.180 (0.034)|
|L1551IRS5||0.85 (0.07)||0.233 (0.037)|
|RNO43||0.81 (0.08)||0.121 (0.034)|
The formulae which connect the integrated intensity of a rotational transition with the number of emitting molecules are, for an optically thin transition, neglecting background radiation;
For a line which has significant optical depth, a correction factor can be applied to the above formulae;
As formaldehyde is abundant in molecular clouds, the H2CO21,1-11,0 transition is likely to have significant optical depth. In order to estimate this we have used the ratios of for the transitions of the main isotopomer and of the 13C substitute.
Combining Eqs. (1) and (2), applying the correction factor, Eq. (4), for the H2CO transition only, assuming a 12C/13C abundance ratio 60, and equal excitation and filling factors for the lines,
then we expect that;
The optical depths for the H2CO transitions, listed in Table 5,
|Source||Optical Depth||Excitation temp. (K)|
Limits on the excitation temperatures of H2CO can be calculated, using the upper limits on the integrated intensity of the H2CO51,4-51,5 transition, via
H2CO and H213CO can exist in both ortho and para forms, depending on the alignment of the spins on the two hydrogens. In this study we have only observed transitions of ortho-H2CO, therefore we need to correct our column densities for these molecules by some assumed ortho/para ratio. The high-temperature statistical value for this ratio is 3:1, however the actual ratio can depend on the temperatures at which the molecules formed, and so may be lower in cold clouds. Kahane et al. (1984), attempted to measure the H2CO ortho:para ratio towards TMC-1. They obtained a best fit to their data of 1:1, but with large associated errors meaning that their observations could also be fit by a higher ratio.
The ortho/para ratio can be used to provide information on the formation mechanisms of molecules. Minh et al. (1995), who observed H213CO in the quiescent cores TMC-1 and L134N, found an ortho/para ratio very close to the statistical value of 3. This suggests that these molecules formed in the gas-phase. Dickens & Irvine (1999) observed H2CO towards star-forming cores, finding the ortho/para ratio to be between 1.5 and 2, indicating that it has been modified due to formation and/or equilibration of H2CO on grains.
We have currently adopted the statistical ratio of 3:1, However we note that adopting a ratio of 1:1 would increase our H2CO column densities by a factor of 3/2, and so reduce the D/H ratios by 2/3, while assuming a ratio of of 2:1, as seen towards other star-forming cores, would only increase the H2CO column densities by a factor of 9/8.
The resulting column densities are given in Table 6.
|Source||5 K||10 K||20 K||30 K||40 K|
|L1448NW||0.065 (0.026)||0.061 (0.013)|
|HH211||0.024 (0.011)||0.022 (0.004)|
|L1527||0.072 (0.021)||0.066 (0.013)|
|L1551IRS5||0.069 (0.017)||0.065 (0.013)|
|HH111||0.031 (0.021)||0.029 (0.014)|
For the HCN1-0 transition we obtained
from the "HFS'' method of CLASS. We then calculated the radiation temperature,
of each triplet by correcting
for the main beam efficiency,
(0.61 at 88.6 GHz) and estimated excitation temperatures,
from the equation;
Column densities for HCN were calculated using;
We calculated column densities for H13CN and DCN assuming optically thin lines (Eqs. (1) and (2)). As we observed two transitions of DCN, excitation temperatures could be calculated via Eq. (6). Values for (DCN) are listed in Table 5.
We note that effects such as beam dilution and/or self reversal in the HCN lines, may lead us to underestimate excitation temperatures for HCN. However, these temperatures are in good agreement, within the uncertainties arising from the spectral noise, with the excitation temperatures of DCN.
Column densities for HCN, H13CN and DCN, along with [DCN]/[HCN] ratios, are given in Table 8.
|(1013 cm-2)||(1012 cm-2)||(1012 cm-3)|
|B5IRS1||2.85 (0.37)||--||1.03 (0.15)||0.036 (0.007)||--|
|L1448mms||8.54 (0.68)||1.26 (0.24)||3.52 (0.39)||0.041 (0.006)||0.047 (0.010)|
|L1448NW||6.77 (0.40)||1.93 (0.29)||4.35 (0.44)||0.064 (0.007)||0.038 (0.007)|
|HH211||4.65 (0.32)||0.85 (0.15)||1.93 (0.25)||0.042 (0.006)||0.038 (0.008)|
|IRAS03282||3.32 (0.30)||0.54 (0.15)||1.37 (0.16)||0.041 (0.006)||0.042 (0.013)|
|L1527||2.51 (0.45)||0.40 (0.10)||0.93 (0.12)||0.037 (0.008)||0.039 (0.011)|
|L1551IRS5||8.42 (0.93)||0.51 (0.10)||1.57 (0.20)||0.019 (0.003)||0.051 (0.012)|
|RNO43||2.02 (0.42)||<0.58||0.75 (0.21)||0.037 (0.013)||>0.022 (0.006)|
L1448NW and L1551IRS5 are the only sources in which [DCN]/[HCN] ratios calculated from observations of HCN do not agree well with those calculated using N(H13CN). This is most likely due to errors in our estimation of the integrated intensity and/or optical depth of the HCN1-0 transitions due to self absorption, and so we prefer the values derived from N(H13CN).
The 31,2-22,1 transition of c-C3H2 has previously been observed towards most of the sources in our survey (Buckle & Fuller 2001) and minimum column densities calculated. We have tentative detections of the 22,0-11,1 transition of the deuterated counterpart of this molecule, c-C3HD, towards 4 sources, B5IRS1, L1448mms, L1448NW and L1527. Column densities and D/H ratios have been calculated, assuming optically thin transitions and an excitation temperature of 10 K, and are listed in Table 9. The upper limits on the [C3HD]/[C3H2] ratios are consistent with observations made in L134N and TMC-1 (Bell et al. 1988), which found [C3HD]/[C3H .
Our values for [HDCO]/[H2CO] and [DCN]/[HCN] are summarised in Fig. 3.
|Figure 3: A summary of our observed a) [HDCO]/[H2CO] ratios (using N(H212CO)) and b) [DCN]/[HCN] ratios (using N(H13CN) (solid error bars) and N(H12CN) (dashed error bars)), plotted against the bolometric temperature of the source, along with least-squares fits to the data (see text), arrows indicate upper limits.|
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To determine the average fractionation, and test if fractionation is a function of evolutionary stage, we have calculated the least squares fits of a straight line to the data in Fig. 3. For the [HDCO]/[H2CO] ratios a good fit could not be found to all the data points, however, when we left out the source HH211, which has a significantly lower HDCO fractionation, we obtained the fit shown as a solid line in Fig. 3a, with uncertainties on this fit shown as dashed lines. This suggests that, with the exception of HH211 and HH111, where [HDCO]/[H2CO] ratios are and , respectively, the observed [HDCO]/[H2CO] ratios are all consistent with a value of 0.05-0.07.
Figure 3b shows the line of best fit to the observed [DCN]/[HCN] ratios (solid line), along with its associated uncertainty (dotted lines). With the exception of HH111, where [DCN]/[HCN] < 0.015, the best fit to our observed ratios is [DCN]/[HCN
which remains flat as
Figure 4a compares the linewidths of H13CN to DCN emission,
|Figure 4: A comparison of the linewidths of the molecules surveyed, the dotted line on each plot indicates equal widths; a) for the 13C-substituted isotopomers of H2CO and HCN vs. for the analogue deuterium-substituted molecules; b) for the main isotopomers vs. for the deuterium-substituted molecules.|
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Our observational results have been interpreted using detailed, chemical kinetic models of deuterium chemistry in dark clouds. The models, which are fully described in Roberts & Millar (2000a,b), now involve 298 species (135 of them containing deuterium) and 5550 reactions.
There are two basic models; the "gas-phase'' model considers only reactions between gaseous species (with the exception that H2 and HD can form on grain surfaces),
|L1448mms||>11.2||1.58 (0.63)||<0.14 (0.06)|
|L1448NW||>9.9||1.59 (0.69)||<0.16 (0.07)|
|L1527||>8.4||2.30 (0.31)||<0.27 (0.04)|
In cold, dense gas, species which collide with dust grains are likely to stick; indeed, it is now well known that interstellar grains become encased in mantles of ice on a similar timescale to the dynamical and chemical evolution of molecular clouds (Willacy & Millar 1998; Rawlings 1999).
The deuterium fractionation, in particular, will be strongly affected. In cold clouds deuterium is extracted from its reservoir in HD by exothermic reactions such as
This theoretical expectation, first predicted by Brown & Millar (1989), has recently been confirmed observationally by the detection of large [DCO+]/[HCO+] ratios in L1544 in clumps in which CO is significantly depleted (Caselli et al. 1999). It may also prove important in explaining the large molecular D/H ratios observed towards IRAS16293 and L134N.
All models presented in this paper use standard depleted solar elemental abundances for dark clouds, as listed in Table 10.
|H2||5.00 10-1||H3+||1.00 10-11|
|He||1.40 10-1||Si||2.00 10-8|
|C+||7.30 10-5||Fe+||1.00 10-8|
|N||2.14 10-5||S||1.00 10-7|
|O||1.76 10-4||HD||3.20 10-5|
Figure 5 shows the variation in [HDCO]/[H2CO] and [DCN]/[HCN] with temperature and density,
|Figure 5: Steady-state [HDCO]/[H2CO] (top) and [DCN]/[HCN] (bottom) ratios, from our gas-phase models, varying over a temperature range 10-50 K (in intervals of 10 K) and a (H2) range of 3.5-6.5 (in intervals of 0.5), the dashed planes indicate the observed [HDCO]/[H2CO] ratios.|
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It appears, therefore, that the gas-phase model alone is insufficient to explain the observations. There are, however, uncertainties inherent in any chemical modelling of this type. In particular, HCN, HNC and CN are all formed via dissociative recombination of HCNH+ with electrons, where our model assumes that 50% of the reactions form CN, while 25% form HCN and 25% form HNC (Herbst 1978). Recent experimental and observational work now suggests that HCN and HNC will be formed in equal amounts, but CN will not be formed by this reaction (Semaniak et al. 2001; Dickens et al. 2000). As there are analogue reactions for forming DCN and DNC, with same rate coefficients and branching ratios, this means that the abundances of HCN, HNC, DCN and DNC would all increase if we adopted these new branching ratios. However, we are primarily interested in abundance ratios, rather than absolute abundances, and as the deuterated and non-deuterated species are being affected in the same way, the molecular D/H ratios will be largely unchanged. We estimate that altering the branching ratios for recombination of HCNH+ (and deuterated analogues) will not alter the predicted [DCN]/[HCN] ratio by more than 50%. It would still be too small to explain all the observations.
The accretion model, on the other hand, produces increased D/H ratios over the gas-phase model, as heavy species freeze out onto grain surfaces. Figure 6 compares results from the gas-phase and accretion models for [HDCO]/[H2CO], [DCN]/[HCN] and other key molecular D/H ratios.
The [DCN]/[HCN] ratio is particularly sensitive to accretion, rising from its steady-state gas-phase value of 0.016 to match the observed ratios of 0.04 in yrs. The DCN abundance is enhanced by constant cycling between DCN and DCNH+, as DCNH+ forms via proton transfer from H3+ and H2D+ to DCN, and via deuteron transfer from H2D+ to HNC. There are similar reactions which form HCN, and so, as the abundance of H2D+ increases, relative to H3+, there is an increased chance of forming DCN rather than HCN.
HDCO and H2CO also react with H2D+ and H3+, forming the ions H3CO+ and H2DCO+ which may then recombine to H2CO and HDCO, the increasing [H2D+]/[H3+] ratio being reflected by the [HDCO]/[H2CO] ratio. However, this is not the only route to forming H2CO and HDCO, therefore, over the same time period the [HDCO]/[H2CO] ratio, is not so sensitive to accretion processes. The ratio rises from 0.05 to 0.063, and so remains consistent with the ratios we measured towards the low-mass cores.
These results assume a kinetic temperature of 10 K and an H2 density of cm-3. [DCN]/[HCN] ratios are sensitive to temperature (see Fig. 5), so if we assumed K, the initial [DCN]/[HCN] ratio would be lower and it would take twice as long (105 yrs) for the predicted ratios to reach the levels we observed. As [HDCO]/[H2CO] ratios increase with temperature for K (see Fig. 5), predicted HDCO fractionation at this time would be much higher than was observed.
Changing the density in the accretion model primarily affects the timescale on which freeze-out occurs; increasing n(H2) from
to 105 cm-3 decreases the time it takes for the predicted [DCN]/[HCN] ratios to rise to agree with the observed values from
to 104 yrs,
|Figure 6: Molecular D/H ratios evolving over time from our gas-phase (left) and accretion (right) models; K, n(H cm-3. The best fits to the [HDCO]/[H2CO] and [DCN]/[HCN] ratios we observed are also indicated on the plots.|
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Table 11 lists the [HDCO]/[H2CO] and [DCN]/[HCN] ratios which have been observed towards other sources.
|IRAS16293||2.0 (-9)1||0.141||1.9 (-9)2||0.012|
|L134N||2.0 (-8)3||0.0684||4.0 (-9)3||0.054|
|TMC-1||7.0 (-8)5||0.0594||1.0 (-8)6||0.0114|
|OCRa||3.7 (-8)7||0.148||2.0 (-8)7||0.049|
|OHCb||2.6 (-8)7||--||3.0 (-7)7||0.0039|
|NOTES: a(-b) implies
a Orion Compact Ridge; b Orion Hot Core.
REFS: 1. Loinard et al. (2000); 2. vanDishoeck et al. (1995);
3. Ohishi et al. (1992); 4. Turner (2001); 5. Ohishi (1998);
6. Willacy & Millar (1998); 7. Charnley et al. (1992); 8. Turner (1990);
9. Schilke et al. (1992).
Shah & Wootten (2001) have observed [NH2D]/[NH3] and [DCO+]/[HCO+] ratios towards a wide range of low-mass, protostellar cores. Their source list overlaps, to some extent, with ours, the relevant results being included in Table 12.
|L1448mms||0.013 (0.004)2||0.029 (0.011)2|
DCO+ is fractionated directly via deuteron transfer from H2D+, therefore [DCO+]/[HCO+] ratios are expected to be among the highest molecular D/H ratios. This is illustrated in Fig. 6, the steady-state [DCO+]/[HCO+] ratio from the gas-phase model, 0.05, rising to 0.08 after yrs of accretion. The [DCO+]/[HCO+] ratio observed towards L1551IRS5 by Butner et al. (1995) is similar to this value, but most of the [DCO+]/[HCO+] ratios observed by Shah & Wootten towards the sources in our survey are lower even than the gas-phase models predict (see Table 12). However, more detailed observations (Shah 2000), indicate that the DCO+ emission may be arising from a different region to the HCO+, thus, these [DCO+]/[HCO+] ratios are very sensitive to the precise details of the radiative transfer model used. Alternatively, this relatively low DCO+ fractionation may simply be evidence that accretion is not currently occurring in these sources, but occurred during an earlier stage of evolution.
Another source which has been observed to have high [DCN]/[HCN] and [NH2D]/[NH3] ratios (0.04 and 0.06, respectively), but relatively low DCO+ fractionation (0.002) is the Orion Compact Ridge (OCR).
Despite the fact that both are classed as hot cores, there are marked chemical differences between the OCR and the Orion Hot Core (OHC) which have been noted by previous workers. It appears that both contain species which have been evaporated from grain surfaces, however, the ice mantles in the OHC seem to have been ammonia rich and methanol poor, while the OCR mantles were methanol rich but ammonia poor.
Levels of deuterium fractionation also differ between the OHC and the OCR. Most ratios which have been measured towards the OHC are 10-3, similar to other hot cores. However, several of the deuterated species observed towards the OCR, including HDCO and DCN, have levels of fractionation similar to or even larger than is seen towards the cold cloud TMC-1.
If the Hot Core formed earlier in the evolution of the ridge cloud, when the gas was mostly atomic, surface hydrogenation would produce species such as NH3, CH4 and H2O, with deuterium fractionation which reflects the atomic D/H ratio in the accreting gas (Charnley et al. 1992; Hatchell et al. 1999). In the Compact Ridge, forming at a later time when the accreting gas was mostly molecular, accretion over a longer timescale at low temperatures may have increased deuterium fractionation. CO in the gas-phase would then form H2CO and CH3OH on the grain surfaces, and deuterium fractionation would be higher in most species.
On the other hand, the differences between the OHC and the OCR could be due to their forming at different temperatures and/or densities (Caselli et al. 1992). Higher temperatures during the collapse phase of the Hot Core would supress deuterium fractionation, and could also mean that CO did not stay on the grain surfaces long enough for significant amounts of H2CO and CH3OH to form.
Hatchell et al. (1998) conducted a survey of [DCN]/[HCN] ratios towards several HMC's, finding ratios typically (0.9-4.1 , later searches for HDS (Hatchell et al. 1999) put limits on [HDS]/[H2S] ratios of 10-3. These values are very similar, and point either to grain surface processing acting to reduce molecular D/H ratios from cold cloud values, or, as deuterium fractionation is so sensitive to temperature variations, to higher temperatures (30 K) at the time the molecules froze onto the grain surfaces.
This appears to contrast with levels of deuterium fractionation observed towards low-mass star forming cores. van Dishoeck et al. (1995) measured [HDS]/[H2S towards IRAS16293, we have measured [DCN]/[HCN towards other low-mass cores. Both these ratios are higher than those observed towards HMC's or predicted by chemical models.
To form stars, cold, dense clouds, with enhanced D/H ratios, collapse under gravity. Gas phase molecules then freeze onto grain surfaces, possibly undergoing some chemical processing, before a shock or outflow disrupts the grains and evaporates their mantles. The fact that this sequence of events produces such different levels of deuterium fractionation in HMC's and low-mass cores may, as in the case of the OHC and the OCR, point to significant evolutionary differences between these two types of star forming region.
Higher temperatures during the pre-collapse phase of high mass stars has already been suggested as an explanation for the lower molecular D/H ratios in HMC's and for the chemical differences between the Orion Hot Core and Compact Ridge. It might also be the case that the longer collapse timescales associated with low-mass star formation mean that accretion processes have significantly impacted on the chemistry in these regions.
We have measured [HDCO]/[H2CO] and [DCN]/[HCN] ratios towards a selection of Class 0 and Class I low-mass, protostellar cores, towards three different star forming regions.
Least-squares fits to our data gave [HDCO]/[H2CO -0.07, consistent with predictions from a gas-phase model at 10 K. Most [DCN]/[HCN] ratios were 0.04, higher than gas-phase models predict, but similar to ratios observed towards L134N and the Orion Compact Ridge.
There are no significant differences between the three different star forming regions we surveyed, or betwen Class 0 and Class I sources. This implies that deuterium fractionation is not being affected by shocks, mantle evaporation or accretion after the Class 0 stage, but is set earlier in the history of the collapsing cloud.
We have used detailed chemical, kinetic models of deuterium fractionation in order to match our observed [HDCO]/[H2CO] and [DCN]/[HCN] ratios. We find best agreement between model predictions and observations by assuming that accretion of molecules onto grains has impacted on the chemistry over a time period of about 50000 years. This is a measure of the timescale over which the gas was cold and dense before forming a protostar. Although it does depend on the density of the gas, the shortness of this period argues that the evolution to high density and the collapse to form a protostar with an outflow occur very rapidly.
The similarity between the [DCN]/[HCN] ratios observed towards these low-mass protostellar cores and those seen in L134N, along with their marked difference from [DCN]/[HCN] ratios observed in high-mass star forming regions, leads us to suggest that there are significant chemical differences between high and low mass star forming regions.
For both high and low mass star forming regions, qualitatively similar processes are occurring. Cold, dense cores, with enhanced deuterium fractionation, begin to contract under gravity. Gas phase species collide with and stick to dust grains, building up mantles of ice in which the species can be oxidised or reduced. Once these cores form a protostar, which heats or disrupts (via an outflow) its surrounding envelope, the grain mantles are released back into the gas-phase. The chemical composition in these cores after evaporation will, therefore, depend on the composition of the ice mantles.
Recent observations (e.g. van der Tak et al. 2000; Nummelin et al. 2001) suggest that ice composition may varies between different environments. These differences may be due to the precise details of the surface chemistry which is occurring (as yet, not well understood), or they may depend on the degree of processing (e.g. by u.v. radiation or cosmic rays) that the ice undergoes. However, the deuterium fractionation within these ices, is likely to be most sensitive to the temperatures and densities under which they formed.
High mass stars appear to form at higher temperatures and densities so molecular D/H ratios which evaporate from grain surfaces in the hot cores close by are lower than those seen in cold clouds. Low mass stars, on the other hand, form at lower temperatures and over longer timescales which means that accretion can impact on the chemistry leading to levels of deuterium fractionation similar to or higher than in dark clouds.
HR is grateful to PPARC for the award of a studentship. Astrophysics at UMIST is supported by PPARC. The 12 m Telescope is a facility of the National Science Foundation currently operated by the University of Arizona Steward Observatory under a loan agreement with the National Radio Astronomy Observatory.