Faraday Discuss. 1998 109 15»30 Chemical models of hot molecular cores Thomas J. Millar* and Jennifer Hatchell Department of Physics UMIST PO Box 88 Manchester UK M60 1QD We present the results of observational studies of a number of molecular clouds associated with ultracompact HII regions in regions of massive star formation. We derive molecular abundances and show that gas-phase-only chemical models of these regions are incapable of reproducing the observations. It appears that a rich chemistry occurs in the ice matrices frozen onto dust grains at high densities in the collapse of molecular clouds to form stars. The switch-on of a protostar can evaporate the ice matrices returning processed material to the gas phase. We present evidence that the surface chemistry can fractionate molecules in deuterium as well as form complex organic molecules.n\1»5. Speci–cally HMCs contain very large abundances of hydrogenated and together with larger species such as and Interstellar clouds in regions of massive star formation are known to contain very small (O0.1 pc; 3]1015 m) hot (T ca. 100»300 K) and dense [n(H2) ca. 1012»1014 m~3] clumps of gas which typically contain several tens of solar masses of gas and have large extinctions towards external radiation. In recent years molecular line observations have shown that these clumps known as hot molecular cores (HMCs) have a diÜerent chemical composition compared to that in cold (10 K) dense [n(H ca. 1010 m~3] dust clouds in which highly unsaturated species are found for example the cyanopolyyne chains 2) HC molecules including CH 2n`1N 4 NH3 H2O H2S CH OH C2H5OH (CH3)2O with abundances enhanced by factors of ca.103»105 over those in cold clouds. Very few unsaturated hydrocarbon chains are detected in 3 HMCs an exception being HC3N which is also observed to be vibrationally excited. Such a diÜerent composition may re—ect the temperature diÜerences between HMCs and cold clouds. However it appears impossible that gas-phase chemistry at 100»300 K is responsible for the bulk of the species observed in HMCs. This is because deuterium fractionation is observed in many molecules including HDO NH2D DCN HDCO and D CO with enhancements of 100»1000 over the cosmic D H ratio of 10~5 similar to those observed in cold clouds.Since the process of fractionation depends on small zero- 2 point energy diÜerences the observed fractionation in HMCs must have taken place at low temperatures (\20 K) and not at the gas kinetic temperature of the HMCs. Taken together with the enhancements in the abundances of hydrogenated molecules the implication is that signi–cant molecular processing has taken place on grain surfaces. Following this reasoning HMCs can be placed into an evolutionary picture of star formation. Cold and dense clouds are in an early or pre-collapse phase. Typical lifetimes before such clouds develop IR sources thought to be indicative of young stellar objects (YSOs) are around a few 105 years. At 10 K the time scale for accretion of gas onto the surfaces of dust grains is about 3]1015/n years where n is the hydrogen number density per m3.Thus on the time scale of cloud collapse and star formation material is expected to collide with and freeze onto the surfaces of cold dust grains. Note that at high number densities the accretion time scale is shorter than the collapse time scale so that all gas 15 16 Chemical models of hot molecular cores phase species with the exception of H2 He and their associated ions eventually –nd themselves incorporated into a frozen ice-like matrix surrounding a refractory grain core. These molecular ices can be detected through IR absorption spectroscopy once the YSO forms but there will be a period before this when both the gas and solid state material become ì invisible œ since gas phase molecules are frozen out and there is no YSO to provide an IR continuum against which to detect absorption.Once massive star formation occurs a hot OB star or more likely a cluster of OB stars will cause thermal heating or drive shock fronts into surrounding dense clumps thereby heating them and releasing the fragile ices into what is now a hot dense and essentially neutral gas. Since grain surface processes tend to produce stable neutrals (radicals are difficult to preserve in the ice phase which may last for 105»106 yr) the resulting clump of gas now an HMC takes over 104 yr to be altered chemically. Eventually parent species evaporated from grain ices are destroyed producing reactive daughter products which can drive a complex chemistry. In this article we shall brie—y review some recent observations of HMCs before discussing detailed chemical kinetic models.We shall show that sulfur species may be used as a chemical clock that the DCN HCN ratio may be used to derive information on the reaction between H and DCN and that models of HMCs need to consider a detailed solid-state chemistry which must be capable of producing molecules as complex as ethanol and dimethyl ether. and Observational results In this section we report on two observational programmes aimed at elucidating the chemical composition and physical condition in HMCs associated with ultracompact HII regions (UCHIIs) that is molecular gas associated with young bright hot stars able to ionise part of their natal molecular cloud.Physical conditions and hot core ages We used the James Clerk Maxwell Telescope (JCMT) to perform a molecular line survey toward 14 HMCs in up to ten 1 GHz bands between 200 and 350 GHz.1 In total some 150 molecular lines have been detected with great variation between sources. Some sources such as G34.3 G10.47 and G31.41 show emission from high excitation transitions of CH OH CH3CN CH3CCH and as well as transitions of complex species such 3 as C2H5OH (CH3)2O C and 2H5CN. Other sources such as G10.30 G13.87 G43.89 and G45.47 contain only a few lines from simple species such as C17O C18O C34S SO and low energy lines of CH Molecular abundances and excitation conditions have been derived using the rota- 3OH CH3CCH. 3OH and CH3CN are thermalized implies number n(H2) in excess of 1013 m~3 which can be reached in the centres of these and CH OH our observations of the 3 tion diagram approach2 which assumes that all emission lines are optically thin that the emission –lls the telescope beam and that the source is in thermal equilibrium at a single temperature.Such diagrams can be made when multiple transitions are observed for a particular molecule. In addition model line intensities can be calculated by adopting local thermodynamic equilibrium (LTE) and varying column density N kinetic temperature T and source size hS . In this case we can treat optically thick transitions. The requirement that the lines of CH densities HMCs. For several molecules such as CH isotopomers indicate that the main lines have signi–cant optical depth so that our LTE 3CN approach is to be preferred.Symmetric top molecules such as CH CN are particularly good tracers of temperature since rotational transitions are split by the K quantum number such that 3 several K components fall close together in frequency thereby minimising observational T . J. Millar and J. Hatchell temperature and source size (in arcsec). errors. We can use our observations of the J\19»18 transition of CH3CN at 349 GHz and the J\13»12 transition at 239 GHz to derive temperatures source sizes virial masses and lower limits to the beam-averaged fractional abundance of CH3CN (Table 1). Since the excitation energy of the observed lines range from 60 to 430 K emission is dominated by hot gas. The actual abundance of CH3CN must be much larger than given in this table because the lines have signi–cant optical depth and the source size is much smaller than the telescope beam.In the strongest sources of CH3CN emission the fractional abundance could be ca. 10~7 approximately 100 times larger than observed in cold dust clouds. Note that although the analysis assumes a constant temperature it appears that a temperature and/or density gradient exists within at least two sources since the J\19»18 transition requires a smaller source size and larger temperature than the J\13»12 transition. Similar results are obtained from analysis of CH OH although the derived temperatures are somewhat lower and the source sizes somewhat 3 larger than those derived from CH CN consistent with the excitation energies of the observed CH 3 3OH transitions and temperature gradients in these sources.Fig. 1 gives CH representative spectra of CH3CN toward G31.41 and Fig. 2 shows model –ts to the 3CN Several HMCs show no CH3CN emission and we derive upper limits to the column density of ca. 1017 m~2 for a temperature of 50 K. SO C34S SO and H2S lines have been detected in a number of HMCs in the 2 2 survey. The 220»211 line of H2S at 216.7 GHz the only accessible transition was detected in all sources observed at this frequency. SO was detected towards six sources with excitation ranging from 100 to 200 K while SO and C34S were detected toward nearly all sources observed. Table 1 CH3CN excitation temperatures source sizes virial masses and beamaveraged fractional abundances (x) in HMCs T /K h /arcsec S source 19»18 13»12 19»18 13»12 Mvir/solar masses 0.98 0.80 0.60 0.63 1.41 1.19 2.29 1.12 1.53 3.46 65 134 141 142 141 119 87 114 149 79 G9.62 G10.47 G29.96 G31.41 G34.3 J\19»18 spectrum with Gaussian –ts to the K\0 1 2 5 and 6 components.Fig. 1 CH3CN spectra towards G31.41. The observed spectra (thin lines) are overlaid with Gaussian –ts to the lines. (a) J\13»12 spectrum with Gaussian –ts to K\0 to 6 components. (b) 17 x]1011 [2 [ » 4 190 810 330 450 460 [7 [3 18 Fig. 2 CH3CN Although our sources were chosen because they are known to emit3 in high excitation lines of NH3 our observations indicate that they can be divided into two broad classes a line-rich class incorporating eight sources which show varying amounts of emission from high excitation lines with G34.3 G10.47 and G31.41 the strongest and G12.21 and G75.78 the least strong of these and a line-poor class which contains six sources.There appears to be a real physical diÜerence between these classes in that the line-rich sources show a hot dense core surrounded by cooler gas while the line-poor sources have evidence only for the cooler halo gas. The hot cores have angular sizes of the order of a few arcseconds corresponding to O0.1 pc temperatures [100 K densities [1014 m~3 and virial masses of the order of a few hundred solar masses. The halo gas is typically cool 10»20 K and dense ca.1010»1011 m~3. The hot cores appear to be centrally condensed in these objects as observations made at oÜsets of 20 arcsec show no emission from hot gas. In addition density and temperature gradients probably exist as can be inferred from the CH3CN observations as discussed above. Toward the cores we derive temperatures of ca. 50 K from CH OH gives temperatures ca. 100 K with increasing tem- while on smaller size scales CH 3 perature and smaller sizes for high excitation transitions. Vibrationally excited CH (and HCN) has also been detected. Emission from these species is pumped by IR absorp- 3CN Chemical models of hot molecular cores temperature and source size for several hot core sources (error bars are 1p) 3CN 19 T .J. Millar and J. Hatchell Fig. 3 The evolution of column densities of various species in the ultracompact component of G34.3 tion of photons emitted by hot dust grains at several hundred Kelvin. Such pumping is efficient over very small angular extents and future observations of vibrationally excited lines particularly with interferometers will be able to trace the hottest regions and hence presumably the source of excitation is these cores. The halo gas is very extended on the order of several parsecs in many cases and has a high column density ca. 1027» 1028 m~2. In the line-poor sources the halo gas appears to have a column density at least one order of magnitude smaller than in the line-rich sources. The lack of detectable hot cores in these sources may be due to beam dilution that is the hot gas has a small angular extent relative to the telescope beam size in which case the hot core sizes would either need to be \0.8 arcsec or that the original hot gas has cooled thus making these sources more evolved than the line-rich sources.Of course another possibility is that they are signi–cantly displaced by at least 10 arcsec from the UCHII regions in which case they would have been missed by our beam. We can also get an estimate of the evolutionary status of the sources through chemical modelling. Charnley4 has suggested that the relative abundances of the sulfurbearing molecules can be used to determine the time from which ice mantles are evaporated. In this picture H2S is formed in the mantles (it is known to have a very low fractional abundance in cold clouds ca.10~11) and evaporates. In addition to proton transfer reactions with H2O` it reacts with atomic H rapidly in hot gas H2S]H]HS]H2 HS]H]S]H2 where the –rst of these reactions has a relatively low activation energy barrier,5 ca. 350 20 Chemical models of hot molecular cores K. The atomic sulfur released by this then reacts with OH and O to form SO which is 2 itself converted to SO and CS. Thus one expects the abundances of H2S SO and SO 2 to trace increasing chemical evolution of evaporated materials from ice mantles. Fig. 3 2 shows the evolution of these abundances in the ultracompact core of G34.3 adopting a temperature of 300 K as a function of time since evaporation of the mantles. At such a high temperature H2S is destroyed completely within a few hundred years a very short time in astronomical terms so that one would never expect to detect any H2S emission from this region.Further out from the centre of the cloud the temperature decreases is this gas at around 100»150 K which gives us the most useful handle on chemical age. and H2S survives longer and contributes appreciably to the observed column density. It factor of 10 between the line-rich sources strongly suggesting that they are at similar In fact the relative abundances of H2S H2CS OCS and CS vary by less than a evolutionary stages. SO and SO abundances are somewhat larger in G5.89 G10.47 and G75.78 than in other sources. These sources may be slightly older or alternatively the 2 excess SO and SO emission may originate from molecular out—ows in or near hot 2 cores.The simultaneous detection of H2S SO and SO places bounds on the source 2 ages of a few thousand to 105 yr depending on the cosmic ray ionization rate and core temperature. Deuterated molecules in hot cores Table 2 presents a list of the deuterated species detected towards three sources including the hot core source in the Orion molecular cloud (OMC) and shows that fractionation is large despite the high kinetic temperature. We have recently begun a programme to study D H ratios in our HMC sample and in May 1997 detected the 3»2 transitions of HC15N and DCN in a number of sources with the JCMT. We also search for emission 20 arcsec oÜ the peak in order to (i) see if the emission was extended and (ii) to derive DCN HCN ratios in the halo gas.Choosing a line of HC15N ensures that emission is optically thin and allows us to directly derive the abundance ratio6 N N (HCN) (DCN)\ [ [ 14 15 N] N] I I D e3.93@Tx x H where [15N]/[14N] is the isotopic ratio of 15N to 14N and varies from source to source,7 but is always within 20% of 400 ID IH are the integrated intensities of the DCN and HC15N transitions respectively and T is the excitation temperature although one can see that the DCN HCN ratio is only weakly dependent on T x x . Table 3 gives the ratio for Tx\30 K. Ratios for Tx[100 K are about 10% larger. The absolute column densities of HC15N and DCN are of course much more sensitive to Tx a parameter we know poorly since we only have one transition observed.Gibb8 estimates T ca. 55 K on core for G34.3 decreasing to ca. 10»20 K at an oÜset of 15 arcsec. Our DCN HCN abundance ratios (1»4)]10~3 are consistent from source to source and an order of magnitude smaller than that observed in the cold dust cloud TMC-1 (0.023^0.001)9 and the Orion ridge cloud (0.02^0.01),10 and similar to the Orion hot core (0.003),10 while our oÜ-core abundance ratios are somewhat larger since we have detected DCN but not HC15N in some positions. It is known that the deuteration ratios in TMC-1 at 10 K and in the Orion ridge at 70 K are dominated by gas phase processes,11 whereas the Orion Hot Core and the HMCs in our survey contain material processed through a grain surface.The observation of DCN HCN ratios as low as 10~3 in these could arise for a number of reasons. Firstly it might re—ect the fact that the ratio was low when the gas accreted onto the ice. A ratio ca. 10~3 requires gas temperatures [50 K for which accretion is inefficient due to thermal desorption.11 T . J. Millar and J. Hatchell Table 2 Fractionation ratios for molecules in TMC-1 (10 K) the ridge clouds of OMC-1 (50»70 K) and the Orion hot core (150»200 K) molecule HDO DCO` N2D` DCN DNC NH C2D HDCO 2D DC C c-C HDCS 3HD 4D DC CH CH 3 5 N N 3 2 OD DOH 2 CH DCCH D2CO Secondly one could argue that a large ca. 10~2 DCN HCN ratio in cold gas could be diluted when frozen onto an ice surface.However detailed models of hydrogenation and deuteration on grain surfaces show that the formation of D-bearing molecules is enhanced over H-bearing molecules partly re—ecting the stronger bonds which D atoms form but also re—ecting the fact that the D H ratio on the surface is greatly enhanced over the cosmic value because it is enhanced in the gas phase and transfers this via collisions on the grain surface. Thus surface chemistry should enhance the DCN HCN abundance ratio.12h15 Finally it is possible that the DCN HCN ratio is large ca. 10~2 or greater on the ice mantles but gets altered on a rapid timescale in hot gas upon evaporation of the ice. Schilke et al.10 suggested that the reaction H]DCN]HCN]D can cycle deuterium between atomic D and DCN and estimated a rate coefficient of 10~16 m3 s~1 and an activation energy barrier E tion barrier the reaction with atomic H alters the initial DCN HCN ratio on a time- A/kB200 K.With such a small activascale of a few hundred years16 for a gas temperature of 100 K. If EA/kT [10 then the timescale for altering the ratio becomes longer than 104 yr incompatible with the ages Table 3 Fractionation ratios (]103) for DCN HCN in various HMCs for an excitation temperature of 30 K source G5.89 G9.62 G10.47 G29.96 G31.41 G34.3 hot core OMC-1 TMC-1 [0.002 0.002 0.003 0.02 0.01 0.045 0.015 \0.045 0.023 0.015 0.01 0.003 0.14 0.02 \0.02 0.015 0.02 0.08 0.004 0.06 0.013 0.03 0.04 0.054 \0.05 0.02 DCN HCN 2.1 2.2 3.2 2.7 2.3 1.4 21 22 Chemical models of hot molecular cores inferred for HMCs.This reaction is endothermic by about 800 K not large enough to prevent H atoms from destroying DCN in hot cores. However recent calculations of the potential surface for this reaction by Herbst and Talbi (this meeting) indicate that energy barriers of several thousand K are present. If this conclusion is correct reactions with H atoms cannot alter the DCN:HCN ratio. NH and Chemical kinetic models of HMCs The earliest models were developed by Brown et al.17 who considered a three stage process. In the –rst stage they followed the chemistry and mantle accretion of a spherical molecular cloud with initial atomic conditions as it underwent a gravitational collapse to high density .In this stage the accretion time was shorter than the collapse time (note taccPn~1 and tcollPn~1@2 where n is the total number density) so that the accreted grain mantle was rich in atoms of C N and O which were hydrogenated to form CH NH 4 and H2O in the second stage which followed a limited surface chemistry consisting essentially of hydrogenation of atomic species. Because of the high mobility of atomic 3 hydrogen (an H-atom scans the entire grain surface in about 10~6 s) hydrogenation is always rapid providing there is a sufficient —ux of H atoms to the surface which can be met at low number densities early in the collapse. At higher densities and later times the conversion of H atoms to H reduces their number and allows heavy atoms to combine.Finally the third stage investigated the chemistry of the evaporated ices in hot gas and 2 found that they survived for at least 104 yr. Brown and Millar13,14 extended this work to include deuteration. 2CO and in ices there is a region in which the hot gas once mantles have been evapo- 3 2CO CH3OH which can be found 3 breakdown leads to a rich N 3 Subsequently much work has concentrated on the third stage4,18h23 and shown that the evaporation of simple molecules from ices can lead to the formation of more complex species in the HMC. Most of these models have been investigated for conditions appropriate to the Orion Hot Core a clump of gas which shows large abundances of N-bearing molecules in contrast to a nearby HMC the Orion Compact Ridge which is rich in O-bearing molecules.The compositional diÜerences between these HMCs are puzzling. Caselli et al.20 suggested an answer through modelling the collapse of a molecular cloud into a central protostar with the inclusion of a radial temperature pro–le. Because of thermal desorption and the diÜerent binding energies involved the composition of the ice mantle varies with radial distance from the young star. In particular the low binding energy of CO means that it is easily evaporated close to the star whereas the much stronger bound NH can be retained. Since CO is hydrogenated to H CH rated contains NH and CO but little or no H 3OH further from the protostar.Because CO is unreactive chemistry close to the protostar in warmer gas while further from the star in relatively cooler gas and H2CO CH3OH can drive the formation of more complex molecules such as (CH3)2O and HCOOCH3 . cores. This is exactly the situation observed in the Orion hot The model developed by Caselli et al.20 used a reasonable description of the grain surface chemistry although the number of reactions included is an order of magnitude less than the number of gas phase reactions. In part this re—ects uncertainty about surface mobilities evaporation rates [which are exponentially dependent on the (uncertain) grain temperature] radical population H-atom —ux reactivities activation barriers and so on. Nonetheless for well observed sources one might hope to build a large enough model to show where gaps in knowledge are crucial.One such source is the HMC associated with the ultracompact HII region G34.3]0.15. Molecular line observations show that the molecular cloud has three main components (i) an ultracompact core, Table 4 Column densities calculated at 3.16]103 yr in the ultracompact core (UCC) and compact core (CC) components at 105 years in the halo UCC species H O CO OH2 NO 2.632]1018 2.700]1019 2.961]1016 1.070]1014 1.724]1016 6.059]1018 4 H CO2O CH2 HNC 2.427]1015 1.129]1017 2.658]1016 6.426]1012 1.196]1016 2.266]1017 HCO HC 2CO 2H2 C C3H2 C4H HC4 C 3N 3H4 1.156]1014 6.853]1014 5.399]1012 2.437]1015 5.573]1013 3.952]1011 He` N CN2 1.131]1019 5.089]1011 6.340]1015 8.181]1014 6.134]1013 5.138]1011 C C3N 3O C CH `2CO NH HCN C2H 2.707]1015 2.880]1016 8.348]1013 4.164]1018 2.254]1014 2.334]1017 NH 3CN CH3 CH3OH H3 ` HCO` N2H` 3.684]1012 8.138]1012 1.184]1012 5.216]1014 2.019]1014 1.453]1011 C2S S C ` 3 S HS H HCS` 2S 1.780]1013 8.942]1012 8.348]1010 3.268]1015 2.582]1017 4.555]1016 2 CS SO SO OCS H HCS 2CS C4S 2.716]1015 2.325]1015 1.542]1014 5.789]1013 2.715]1015 4.845]1013 C NS ONS Si CH CHO 3.077]1016 1.115]1017 2.242]1017 1.942]1016 1.682]1014 3 T .J. Millar and J.Hatchell halo CC 2.659]1020 1.646]1019 2.333]1018 1.336]1016 6.973]1016 2.210]1017 8.567]1018 1.346]1019 2.783]1017 1.877]1016 5.831]1016 2.468]1018 9.237]1015 6.257]1016 8.802]1014 6.583]1012 3.483]1015 7.567]1015 2.626]1016 5.804]1016 9.851]1015 4.996]1012 1.235]1016 1.170]1017 1.413]1015 5.261]1016 6.742]1014 1.657]1014 1.002]1013 4.208]1014 1.280]1015 3.184]1014 3.835]1012 1.928]1015 1.657]1013 3.325]1012 5.921]1018 1.855]1015 1.095]1014 6.096]1014 1.091]1014 9.185]1014 5.554]1018 4.846]1013 3.034]1015 8.763]1014 2.109]1013 3.127]1012 1.165]1014 6.373]1014 1.229]1015 1.324]1016 3.753]1013 4.222]1013 4.352]1015 7.299]1015 6.087]1014 2.085]1018 5.045]1013 1.137]1017 2.012]1015 1.151]1015 1.767]1014 1.861]1015 9.452]1014 3.566]1015 3.124]1013 4.925]1013 9.744]1012 1.695]1014 7.768]1013 1.832]1012 4.718]1012 1.785]1013 7.153]1014 8.707]1016 8.700]1016 2.799]1016 6.229]1014 1.716]1017 7.223]1011 1.619]1015 1.361]1016 1.414]1016 1.048]1015 2.089]1015 2.732]1013 3.715]1013 3.491]1012 2.683]1017 2.830]1013 3.899]1014 4.864]1012 5.795]1012 3.504]1015 9.056]1014 1.572]1019 4.651]1018 1.190]1017 1.465]1015 4.679]1012 1.029]1017 2.106]1017 6.394]1016 9.463]1015 9.400]1013 23 observedb totala [2.9]1019 2.771]1020 5.692]1019 2.641]1018 3.224]1016 1.453]1017 8.748]1018 [5.8]1015 1.3]1019 \3.3]1014 [5.3]1013 [2.3]1016 3.792]1016 2.335]1017 3.732]1016 1.801]1013 2.780]1016 3.512]1017 [2.5]1015 1.8]1016 2.808]1015 5.362]1016 6.835]1014 4.530]1015 8.232]1013 4.245]1014 [5.9]1014 6.7]1014 2.279]1019 1.904]1015 9.483]1015 2.304]1015 1.916]1014 9.221]1014 7.176]1015 3.674]1016 1.921]1015 6.262]1018 3.134]1014 3.471]1017 1.2]1015 [8.1]1015 [2.7]1018 2.4]1014 3.7]1016 [2.0]1015 2.047]1015 1.208]1015 1.876]1014 2.552]1015 1.225]1015 3.568]1015 6.454]1014 1.716]1017 7.161]1014 9.196]1016 3.588]1017 8.769]1016 [1.1]1016 [3.0]1013 1.0]1016 [2.4]1015 1.6]1016 [7.7]1015 [5.6]1016 [4.0]1013 3.792]1015 4.804]1015 1.864]1014 1.008]1014 6.223]1015 2.693]1017 1.585]1019 4.973]1018 4.071]1017 3.035]1016 2.668]1014 24 Chemical models of hot molecular cores Table 4»(Continued) halo CC UCC observedb species totala 3.5]1015 C2H5OH CH OCH 3 HCOOCH3 1.6]1016 3 1.175]1010 5.266]1014 2.647]1014 4.623]109 3.569]1015 8.875]1012 3.684]108 1.748]1011 1.006]1011 1.971]109 2.282]1012 1.595]106 1.026]1010 2.956]1014 7.101]1014 1.043]109 1.367]1015 4.236]1012 1.125]109 2.309]1014 1.936]1014 1.609]109 2.200]1015 4.639]1012 C2H6CO C C2H4 2H5 5.6]1023 C2H6 H electron 2 3.510]1015 5.706]1023 7.293]1016 0.000]10 1.657]1023 2.834]1016 1.206]1015 1.350]1023 3.403]1016 2.304]1015 2.700]1023 1.056]1016 a Total column density to the cloud centre.b Column density derived from observation. 2. H2O On the other hand with a size of 0.01 pc n(H2) ca. 2]1013 m~3 T ca. 300 K (ii) a compact core with size n(H 0.1 pc ca. 2]1012 m~3 T ca. 100»300 K and (iii) a molecular halo with size ca. 2) n(H 3.5 pc 2)Pr~2 and T (r)Pr~0.4. Macdonald et al.24 performed a 330»360 GHz spectral line survey using the JCMT and detected some 338 lines from 35 species and estimated molecular abundances. Further observations were done on this source as part of the HMC survey.1 These observations show evidence of all three components in G34.3 and a detailed chemical kinetic model has been constructed.25 The model follows the chemistry at 22 radial points.In the halo evolution from atomic constituents is followed whereas in the compact and ultracompact cores where mantle evaporation occurs the chemistry follows the breakdown of parent species. Except for H formation no grain surface chemistry is followed but the gas phase chemistry is extensive with ca. 2 2200 reactions covering 225 species. Because the chemistry is depth-dependent column densities can be calculated as a function of time. Fig. 3 shows how the column densities of selected species change in the ultracompact core. One sees that H2S is destroyed (by H atoms) within a thousand years to produce SO and SO which possesses a large activation barrier for reaction with H atoms is essentially unchanged.The ice mantles are assumed to be fairly simple and one sees that reasonable column densities of some complex molecules can be found in particular CH CHO HCOOCH 3 and (CH3)2O 3 but no C2H5OH is formed whereas the observational abundance is26 ca. 10~8»10~7 typical of many HMCs. This failure to reproduce the C2H5OH abundance may re—ect a lack of knowledge about its gas phase production in hot gas or more interestingly be an indication that C2H5OH is actually a parent molecule that is a product of grain surface chemistry itself. and The detailed model of G34.3 shows that good agreement can be found for most molecules if the time scale from evaporation is ca. 3000»104 yr and that for evolution of the halo gas is ca.105 yr (see Table 4 for a comparison for calculated column densities with those observed) but clearly points to species such as C2H5OH and HCOOCH3 which are ca. 105 and 102 short of the column densities observed. For this reason we have been developing a comprehensive time dependent chemical kinetic model which follows both the gas phase and grain surface chemistry in a collapsing cloud with a grain surface chemistry detailed enough to account for the synthesis of C2H5OH (CH3)2O.27,28 The need for surface chemistry to produce complex species can be in a number of HMCs. 2H5OH (CH3)2O inferred from Table 5 which gives observed fractional abundances29 for CH OH and C 3 In our model we have included 80 surface species and 200 gas phase species limited and (CH by 2000 gas phase reactions and about 80 surface reactions based on those described by Hasegawa and co-workers.30,31.However our surface reactions are optimised to allow the production of CH OH C2H5OH 3)2O on grains. In addition to simple 3 25 T . J. Millar and J. Hatchell hydrides we include the surface formation of CHO and H2CO CO2 by reactions such as *CO]*H]*HCO *HCO]*H]*CO]H2 *CH2]*O]*H2CO *H2CO]*H]*HCO]H2 2]*H *CO]*OH]*CO *HCO]*O]*CO2]*H where *M represents a chemical species M on the grain. However important reactions which mitigate against the formation of molecular backbones containing the CwC bond have been neglected in particular *C]*C]*C2 which could eventually lead to the production of ethane *C2H6 which is detected in cometary comae.32 Because of this neglect our models calculate the maximum abundances of complex molecules possible.The majority of reactions which build alcohols and ethers involve the reaction of *CH (n\0»3) with *OH (m\0 1) m n 3 *CHn]*OHm ]*CHnOHm OH. Routes to ethanol and followed by hydrogenation as necessary to produce *CH dimethyl ether involve forming a backbone of either CwCwO or CwOwC followed by hydrogenation. The backbones form via *CO]*C]*CCO *CO]*C]*COC A large number of model calculations have been performed. The models consider freefall gravitational collapse from initial densities n of 108 109 and 1010 m~3 as well as i retarded collapse models in which the collapse is slowed by a factor B. Free-fall times range from 4.4]106 yr for ni\108 m~3 to 4.4]105 yr for ni\1010 m~3.Changes in visual extinction and temperature are also followed during the collapse. One should note that in the late stages of collapse molecular line emission becomes optically thick and the cloud is not able to radiate away its gravitational energy. As a result it heats up T such that once dust[50 K mantles evaporate rapidly. In addition to this thermal desorption process the model also contains a number of other desorption mechanisms including mantle explosions direct UV photodesorption cosmic ray induced photodesorption and direct heating by cosmic rays. In addition some models not presented here consider desorption via periodic shock waves and thermal desorption caused by material in a rotating cloud passing close to a nearby star.Table 5 Observed fractional abundances of complex organic molecules in HMCs source C2H5OH CH3OH (CH3)2O SgrB2(N) OrionCR W51e1/e2 G34.3]0.15 NGC6334F 1.3]10~9 3.0]10~9 9.5]10~9 1.0]10~9 2.4]10~9 4.3]10~8 9.3]10~7 9.5]10~7 9.4]10~8 4.8]10~7 2.0]10~8 1.3]10~6 1.0]10~7 3.9]10~7 8.8]10~7 26 Chemical models of hot molecular cores Fig. 4 The evolution of fractional abundances of various gas phase species in the –rst 2]106 yr of a retarded collapse Fig. 4»6 show sample gas phase and grain surface fractional abundances for a retarded collapse from 109 m~3. In Fig. 4 the gas phase abundances are shown for times up to 2]106 yr. At times bigger than this accretion of the gas begins to dominate and abundances decrease (Fig.5) essentially to zero. Correspondingly in this period ice 27 T . J. Millar and J. Hatchell Fig. 5 The evolution of fractional abundances of various gas phase species between 2 and 4]106 yr in a retarded collapse 2H5OH\10~8 and of (CH3)2O ca. 10~6 (Fig. 5). mantles build up and surface chemistry builds complex species (Fig. 6). For t[4]106 yr thermal desorption liberates the mantles and returns material to the gas phase giving fractional abundances of C The calculations explicitly follow the gas and surface abundances and the exchange of material as a function of time so one can also obtain detailed information on the depth dependence of the composition of the ice mantle which in this case amounts to 28 Chemical models of hot molecular cores some 130 layers of material.Fig. 7 shows the percentage contribution to each layer for a variety of species as well as the variation of number of layers as a function of time. The dramatic increase of the number of layers at late times is a direct consequence of the retarded free-fall collapse for which the timescale to collapse to an in–nite density is proportional to n~1@2. Thus at low number densities the cloud evolves slowly. Once it reaches high density it both accretes material rapidly and collapses quickly. At early times the mantle is mainly composed of H2O since O atoms are the most abundant abundant CO is formed in the gas and accreted. The fractional abundance of *CH OH species colliding with the grains.At late times the contribution of H2O decreases as *C 3 2H5OH *(CH3)2O and together with their percentage contribution to the mantle are given in Table 6 for the model considered here. We note that on evaporation this would provide fractional abundances of CH OH C2H5OH and (CH3)2O 3 compatible with those observed (Table 5). in the gas phase Conclusions Our molecular line survey shows that hot molecular cores associated with UCHII regions can be rich sources of molecular emission with densities and temperatures much larger than is typical for molecular clouds. The chemical composition of the hot gas re—ects both the direct and indirect products of evaporated mantle ices and we have been able to use observations of sulfur-bearing molecules to constrain the ages of Fig.6 The evolution of fractional abundances of various grain mantle species between 2 and 4]106 yr of a retarded collapse. For times greater than 4]106 yr the mantles have evaporated. 29 T . J. Millar and J. Hatchell Fig. 7 The evolution of the number of mantle layers and the percentage contribution of various species to the mantle composition as a function of mantle layer number HMCs. In addition our observations of the DCN HCN abundance ratio shows evidence for the importance of the destruction of DCN by atomic hydrogen and allows inferences to be drawn about the likely activation energy barrier in this reaction. We have also reported on the results of detailed chemical kinetic models of hot cores and of the gas»grain interaction and surface chemistry in collapsing cores.Although there are several uncertainties in this latter model including the details of the surface chemistry in particular the possible presence of activation energy barriers stoichiometric eÜects binding energies and mobilities in realistic ices the neglect of a more comprehensive reaction network and the sensitivity to the details of the collapse since this determines the abundances of gas phase species which collide with the grains it is 30 *(CH3)2O % time/yr abundance 0.006 1.05]10~8 4.16]106 Chemical models of hot molecular cores Table 6 Calculated fractional abundances and percentage contribution of complex organic molecules in the ice mantle at the time of maximum mantle depth *C2H5OH *CH3OH % abundance % abundance 0.006 4.02 1.00]10~8 6.74]10~6 encouraging that a set of models does emerge for which reasonable agreement between both surface abundances when compared to those observed through IR absorption of protostellar sources and cometary ices which appear to be very similar to interstellar ices can be found.It appears that surface chemistry in cold interstellar ices is capable of producing a range of complex organic molecules. 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