19. Interstellar Chemistry

Transcription

1 19. Interstellar Chemistry 1. Introduction to Interstellar Chemistry 2. Chemical Processes & Models 3. Formation & Destruction of H 2 4. Formation & Destruction of CO References Duley & Williams, "Interstellar Chemistry" (Academic 1984) van Dishoeck, UCB Lectures Notes (2000) Genzel, Sec. 3 of "Physics & Chemistry of Molecular Clouds" in "Galactic & Interstellar Medium" (Springer 1992) van Dishoeck & Blake, ARAA, 36, 317, 1998 van Dishoeck, ARAA, Data Collections V. Anicich, ApJS 84, 245,

2 1. Introduction to Interstellar Chemistry Astrochemistry deals with the formation and destruction of multi-atomic systems of interest to astronomy and planetary science, including stellar and planetary atmospheres and comets. The systems of interest range from diatomic molecules like H 2 and CO through large molecules like PAHS, dust grains and even larger bodies. We focus on the new field of interstellar chemistry whose growth was stimulated by the discovery of molecules in the ISM and which also encompasses nano-through micronsized particles. Two basic differences between interstellar chemistry and terrestrial chemistry are: 1. Much lower pressures 2. Extreme temperatures from cryogenic to ultra-hot 3. External energy sources, e.g., far-uv, CRs, etc.

3 Special Features of Interstellar Chemistry Cosmic Rays Dynamical Changes UV Photons Cosmic Mix of Elements n,t,x Each gas element is an open system subject to external effects. The low ISM pressure means that the main phases are gas and solid, bridged by large molecules & solids. The chemistry is heterogeneous, with reaction between gas and solid, especially in the surface layers of solids. The chemistry is often time-dependent, i.e., chemical timescales may be long compared to thermal & dynamical.

4 The Importance of Interstellar Chemistry More than 140 molecules have been detected (not counting isotopic variants). Examination of the list (next page) suggests some generalizations: Small chains favored over rings. Many unsaturated radicals (20%) and ions (10%). Some astrophysical objectives of interstellar chemistry: (1) Model chemical observations in terms of temperature, density, & ionization fraction (2) Find diagnostics for total elemental abundances and total hydrogen density, e.g., relate measurements of OI to the total abundance of oxygen or CO to H 2.. More generally, find diagnostics of the physical parameters.

5 The Interstellar Molecules

6 Types of Chemical Models Shielded dense cloud Partially transparent diffuse clouds Photon-dominated regions Collapsing clouds Shocks Stellar winds Cloud cores & YSO envelopes Disks Modeling interstellar chemistry became intense after much of it Influenced by Salpeter & Co. (e.g., Gould, Werner, Watson, & Draine). Other pioneering work emphasizing grain catalysis was by Bates & Spitzer (ApJ ) & Stecher & Williams (ApJ ). Most chemical modeling after 1973 starts with gas-phase chemistry and later attempts to include surface interactions.

7 2. Gas Phase Chemistry We begin with gas-phase chemistry for cool (T < 100 K) low-ionization (x e < 10-3 ) gas with minimal attention to surface chemistry and primary application to: Diffuse Clouds and Exteriors of Molecular Clouds Moderate density - typically 100 cm -3 Low temperature - typically 50 K Low ionization - mainly H + and C + Molecular Cloud Interiors Intermediate density - typically 500 cm -3 Low temperature - typically 20 K Very low ionization - typically 10-7 NB: all clouds are inhomogeneous.

8 Chemical Activity Most chemical reactions do not proceed under the above conditions, e.g., the most abundant molecules in cool dense clouds, H 2 and CO, do not react on collision. Even exothermic reactions may not proceed unless the kinetic energy is larger than some characteristic activation energy of order 1 ev. To activate cold chemistry, energy has to be provided externally in the form of UV photons, cosmic rays, or mechanical energy. Energizing the chemistry is particularly important in partially unshielded regions where UV photons photodissociate molecules at the rate s -1 (on a timescale as short as 300 yr)

9 1. One-body Generic Reaction Types hν + AB A + B : at a volumetric rate G(AB) n(ab) CR + AB AB + + e : at a volumetric rate ς (AB) n(ab) 2. Two-body A + B C + D at rate per unit volume k (A,B; T ) n(a) n(b) = k(a,b; T ) x(a) x(b) n The rate coefficient is often a function of T, as in reactions with activation energy, where it is proportional to exp(-t a /T). Conclusion: Reaction rates are sensitive to both T and n H. 2 H NB Three-body reactions are rare because of the low density of the ISM.

10 Examples of Specific Reactions Many of these are able to overcome the weakness of ordinary chemical reactions in cool regions Neutral Radical Warm regions (> 300 K); k up to cm 3 s -1. O + H 2 OH + H OH + H 2 H 2 O + H O + OH OH 2 + O Ion-Molecule Important for cool regions; often k ~10-9 cm 3 s -1 in the absence of activation energy (most very exothermic reactions) charge exchange H + + O O + + H (H) abstraction O + + H 2 OH + + H proton transfer H 3+ + O OH + + H 2 Radiative Association - Weak, but may sometimes be the only pathway, e.g., in the dark ages e + H H - + hν H + H + H hν

11 Photodissociation Almost always important; typical ISM rate ~10-10 s -1. hv + H 2 2H hv + OH O + H (via Lyman and Werner band lines) N.B. Many cross sections have not been measured accurately Dissociative Recombination - Always important; k~10-7 cm 3 s -1 e + H 3+ 3H, H 2 + H (branching 3:1) e + H 3 O + OH + H 2, OH + 2H, H 2 O + O Addendum on Surface Reactions Important but poorly understood, e.g., H + H(Gr) H 2 + Gr (geometric rate; e.g., Hollenbach & Salpeter 1971) Additional processes: adsorption & desorption; catalysis; icing!

12 3. Formation and Destruction of H 2 A. Formation in Cool regions (T < K) The simplest process, radiative association of two H atoms, H + H H 2 + hν, is very weak (a ro-vibrational, quadrupole transition). For more than 40 years it has been generally assumed that H 2 is formed on grains, as in volumetric formation rate = R n H n(h) where, following Hollenbach & Salpeter (1971), R is essentially the rate at which H atoms strike the available grain surface area (per unit volume) with a sticking factor of 1/3, R = 3 X cm 3 s -1 (T/100K) 1/2. Recent laboratory experiments by Vidali et al. raise questions on how the recombination actually occurs and lead to a rate coefficient that is significantly smaller (e.g., Lipshat et al. MNRAS ).

13 Formation of H 2 (cont d) B. Formation in Warm Regions ( T > 300 K ) radiative association of H ions of both signs 1. H + + H H 2+ + hv 2. e + H H - + hv H 2+ + H H 2 + H + H - + H H 2 + e Typical applications are post-recombination and outflows from YSOs. C. Formation in Very Dense Regions ( > cm 3 s -1 ) - three-body reactions. e.g., 3H H 2 + H 2H + H 2 2H 2 with rate coefficients cm 6 s -1. Typical applications are cool stellar atmospheres, YSO disks and jets.

14 Photodissociation of H 2 As discussed in Lecture 18, H 2 is dissociated in a two-step process when a far UV photon is absorbed in a Lyman or Werner transition below 1108 Å, and the excited molecule decays into the continuum of the ground electronic state with the emission of a photon in the Å band. With typical oscillator strengths of 0.03, the L-W lines easily become optically thick. For example, with a Doppler parameter b = 3 km s -1, the standard formula for optical depth at line center τ ( 0) = f lu 2 πe m c e π b yields N(H 2 ) > cm -2. Thus absorption occurs mainly in the damping wings of the Lyman and Werner lines. The theory of H 2 photodissociation was developed by collaborators & students of Salpeter, e.g., Hollenbach, Werner, & Salpeter (ApJ,163, 166, 1971), and applied to the early Copernicus observations by Jura (ApJ ) and Federman, Glassgold & Kwan (ApJ, 227, 466, 1979; FGH) N l

15 Photodissociation of H 2 (cont d) FGH showed that, in the presence of H 2 absorbers, the transmitted dissociating radiation is given by the derivative of the equivalent width. The simple proof follows by considering the definition of equivalent width (e.g., Lecture 2) for the transition j k: W [ exp( N s ϕ( ] = dνϕ( ν ) 1 ν )) j Here σ ( ν)=s φ(ν) is the frequency dependent absorption cross section. Differentiation leads to: ( N = W j s dv s ) dvϕ ( ν )exp( Njsϕ( ν )) ϕ( ν )exp( N = dvϕ ( ν ) s j s ϕ( ν )) ϕ( ν ))

16 This ratio is called the line self-shielding function J W ( N s j ) == dv s ϕ( ν )exp( N dvϕ ( ν ) s j s ϕ( ν )) ϕ( ν )) It gives the absorption cross section, averaged over the line shape, taking into account the attenuation of the dissociating radiation by H 2 itself; it can be considered a function of the optical depth at line center J τ (0) = N s ϕ(0) = j f 2 πe m c e N l π b (for Doppler broadening) J is a decreasing function of optical depth (or column), starting from unity at zero optical depth. Having already analyzed the dependence of W on optical depth, FGK obtain the dependence of J by differentiation.

17 Behavior of the Curve of Growth & Line Self-Shielding Function There are three regimes of equivalent width, depending on optical depth at line center: τ 0 << 1 linear ~ τ 0 τ 0 > 1 flat or logarithmic ~ logτ 0 1/2 τ 0 >> 1 square-root ~ τ 0 1/2 Differentiating W gives the corresponding behavior for J: τ 0 << 1 J = 1 τ 0 > 1 J ~ 1/τ 0 τ 0 >> 1 J ~ τ 0-1/2 In practice, the Lyman & Werner bands are usually optically thick, and the dissociation rate decreases as the inverse square root of the H 2 column density

18 Abundance of H2 in Diffuse Clouds (and The Edges of Thick Clouds) Follow FGK (or Warin et al. A&A for a more detailed theory). Balance the photodestruction rate G 0 Jn(H 2 ) against the grain formation rate Rn H n(h) while satisfying mass conservation, n H =n(h) + 2n(H 2 ). The photo rate is calculated summing all accessible Lyman & Werner band transitions. The effects of other absorbers of Å radiation are also easy to include. M = x(h ) = 1+ ε cm, ε = R = 3 10 G 2n 0 H J R cm s ( T J /100K) M N(H / ), G = 6 10 s ε is the ratio of the formation to the photodissociation time scales. Grain formation is a slow process, e.g., the formation time at 100 K is 1 Rn H 100cm = 10 Myr nh -3

19 Spatial Variation of the H-H 2 Transition N(H 2 ) dashed curve: Glassgold & Langer (ApJ , similar to HWS) a. refined FGK theory (ground state) b. with dust c. thermal level population d. collisional-radiative population dotted-curve: 10% conversion N H dots - Copernicus data line - FGK theory Limited observational support for the theory of H 2 formation & destruction

20 4. Formation and Destruction of CO A. Formation in Cool Regions: Fast ion-molecule reactions k ~10-9 cm 3 s -1 O + + H 2 OH + + H H 3+ + O OH + + H 2 OH + + H 2 OH 2+ + H or H 3+ + OH OH 2+ + H 2 OH 2+ + H 2 OH 3+ + H H 3+ + H 2 O OH 3+ + H 2 and fast dissociative recombination, β ~10-9 cm 3 s -1 : OH + + e O + H OH2 + + e OH + H and O + H 2 or 2H OH3 + + e H 2 O + H and OH + H 2 or 2H, etc. 22% OH 75% OH OH is a crucial intermediary in the formation of CO, and it requires H 2 as a precursor. But the synthesis of OH is driven by the ions O + and H 3 + (to be discussed in Lecture 20). This scheme also produces H 2 O and other O-bearing molecules, as will also be discussed.

21 Gas-Phase Synthesis of CO The efficiency of ion-molecule synthesis of OH depends on the ratio of reaction rates representing the competition between formation and destruction, k x(h 2 )/ βx(e). In diffuse clouds with x(e) ~10-4 and x(h 2 ) > 0.01, this ratio is ~ 1. CO is rapidly produced from OH by fast reactions, e.g., C + + OH CO + + H or, where C is more abundant CO + + H 2 HCO + + H than C +, the neutral reaction C + + H 2 O HCO + + H is fast: C + OH CO + H HCO + + e CO + H The idea of using ion-molecule reactions to initiate low-temperature interstellar chemistry is due to WD Watson (ApJ 183 L ) and Herbst & Klemperer (ApJ ), although the latter did not give the correct theory for the synthesis of CO. Klemperer was responsible for the identification of an unknown mm transition with HCO+ (ApJ ).

22 Cosmic Rays and the Synthesis of OH and CO Where do the ions required by ion-molecule chemistry come from? In cool regions, the usual answer is cosmic rays. In Lecture 14, we estimated the cosmic ray ionization rate on the basis of demodulated satellite observations out to 42 solar radii as ς CR ( > 2 MeV) 5 10 This refers to the ionization of atomic H, i.e, CR + H H + + e Similar rates applies to CR ionization of He and H 2, CR + He He + + e CR + H 2 H 2+ + e and H + + H (5%) But H 2+ is very reactive: H H H + + H 2 H 2+ + H 2 H 3+ + H 2 and leads to H 3+, the pivotal ion in the ion-molecule chemistry. 17 s 1

23 Formation and Destruction of CO (cont d) B. Formation in Warm Regions - Radical reactions: O + H 2 --> OH + H C + OH --> CO + H (slow) (moderately fast) Destruction of CO 1. In partially UV-shielded regions, the same two-step process operates as for H 2, with the following changes: a. G 0 = 2x10-10 s -1 (for the mean UV radiation field), b. self-shielding is reduced since usually x(co) << x(h 2 ). c. mutual shielding in the Å band by H 2 and even some lines of atomic H are importnat d. mutual shielding of the isotopes 13 C and C 17 O by the lines of 12 CO and H 2 occurs 2. In well-shielded (thick) clouds, CO is destroyed by the fast ionmolecule reaction He + + CO C + + O + He.

24 Synthesis of CO: Summary All of the ion-molecule reactions for CO are fast, but the ions may be in short supply in dense shielded regions because their abundances depend on ς CR /n H. In partially shielded regions the time scales are short, but only a small fraction of carbon goes into CO because of strong photodissociation. Slow formation of H 2 may be the limiting factor in the synthesis of CO in atomic regions. Much depends on the initial conditions and whether steady state can be reached..