19. Interstellar Chemistry

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, 42 119 2004

Data Collections• V. Anicich, ApJS 84, 245, 1993• www.rate99.co.uk• http://kinetics.nist.gov/index.php

1. Introduction to Interstellar Chemistry2. Chemical Processes & Models3. Formation & Destruction of H24. Formation & Destruction of CO

1. Introduction to Interstellar ChemistryAstrochemistry 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 H2 and CO through large molecules like PAHS, dust grains and even larger bodies.

We focus on the new field of interstellar chemistry whosegrowth was stimulated by the discovery of molecules in theISM and which also encompasses nano-through micron-sized particles. Two basic differences between interstellar chemistry and terrestrial chemistry are:

1. Much lower pressures2. Extreme temperatures from cryogenic to ultra-hot3. External energy sources, e.g., far-UV, CRs, etc.

Cosmic Mix of Elementsn,T,x

Cosmic Rays UV PhotonsDynamical Changes

• Each gas element is an “open” system subject to external effects.

• The low ISM pressure means that the main phases aregas 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.

Special Features of Interstellar Chemistry

The Importance of Interstellar Chemistry

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 H2.. More generally, find diagnostics of the physical parameters.

More than 140 molecules have been detected (not counting isotopic variants).Examination of the list (next page) suggests somegeneralizations: Small chains favored over rings. Many unsaturated radicals (20%) and ions (10%).

The Interstellar Molecules

Types of Chemical ModelsShielded dense cloudPartially transparent diffuse cloudsPhoton-dominated regions Collapsing cloudsShocksStellar winds Cloud cores & YSO envelopesDisks

Modeling interstellar chemistry became intense after 1973. much of it Influenced by Salpeter & Co. (e.g., Gould, Werner, Watson, & Draine). Other pioneering work emphasizing grain catalysis was by Bates &Spitzer (ApJ 113 441 1951) & Stecher & Williams (ApJ 146 88 1966).

Most chemical modeling after 1973 starts with gas-phasechemistry and later attempts to include surface interactions.

2. Gas Phase ChemistryWe begin with gas-phase chemistry for cool (T < 100 K) low-ionization (xe < 10--33) gas with minimal attention to surface chemistry and primary application to:

Diffuse Clouds and Exteriors of Molecular CloudsModerate density - typically 100 cm-3

Low temperature - typically 50 K Low ionization - mainly H+ and C+

Molecular Cloud InteriorsIntermediate density - typically 500 cm-3

Low temperature - typically 20 K Very low ionization - typically 10-7

NB: all clouds are inhomogeneous.

Chemical ActivityMost chemical reactions do not proceed under the aboveconditions, e.g., the most abundant molecules in cool dense clouds, H2 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 photo-dissociate molecules at the rate 10-10s-1 (on a time-scale as short as 300 yr)

Generic Reaction Types

)AB()AB( rate c volumetriaat :eABAB)AB()AB( rate c volumetriaat :BAAB

nCRnGh

ς

ν

+→+

+→++

1. One-body

2. Two-body

2H(B)(A))B;(A,(B)(A))B;(A,

eunit volumper rateat DCBA

nxxTknnTk =

+→+

The rate coefficient is often a function of T, as in reactionswith activation energy, where it is proportional to exp(-Ta/T).Conclusion: Reaction rates are sensitive to both T and nH.

NB Three-body reactions are rare because of the low density of the ISM.

Examples of Specific ReactionsMany of these are able to overcome the weakness of “ordinary”chemical reactions in cool regions

Neutral Radical – Warm regions (> 300 K); k up to 10-10 cm3 s-1.

O + H2 OH + H OH + H2 H2O + H O + OH OH2 + O

Ion-Molecule – Important for cool regions; often k ~10-9 cm3 s-1 in the absence of activation energy (most very exothermic reactions)charge exchange H+ + O O+ + H (H) abstraction O+ + H2 OH+ + H proton transfer H3

+ + O OH+ + H2

Radiative Association - Weak, but may sometimes be the only pathway, e.g., in the dark ages

e + H H- + hν H + H+ H2 + + hν

Photodissociation – Almost always important; typical ISM rate ~10-10 s-1.

hv + H2 2H (via Lyman and Werner band lines)hv + OH O + H

N.B. Many cross sections have not been measured accurately

Dissociative Recombination - Always important; k~10-7 cm3 s-1

e + H3+ 3H, H2 + H (branching 3:1)

e + H3O+ OH + H2, OH + 2H, H2 O + O

Addendum on Surface Reactions – Important but poorly understood, e.g.,

H + H(Gr) H2 + Gr (geometric rate; e.g., Hollenbach & Salpeter 1971)

Additional processes: adsorption & desorption; catalysis; icing!

3. Formation and Destruction of H2A. Formation in Cool regions (T < 100-200 K)

The simplest process, radiative association of two H atoms,H + H H2 + hν,

is very weak (a ro-vibrational, quadrupole transition). For more than 40 years it has been generally assumed that H2 is formed on grains, as in

volumetric formation rate = R nH 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 10-17 cm3 s-1(T/100K)1/2.

Recent laboratory experiments by Vidali et al. raise questions on howthe recombination actually occurs and lead to a rate coefficientthat is significantly smaller (e.g., Lipshat et al. MNRAS 348 1055 2004).

B. Formation in Warm Regions ( T > 300 K ) – radiative associationof H ions of both signs

1. H+ + H H2+ + hv 2. e + H H- + hv

H2+ + H H2 + H+ H- + H H2 + e

Typical applications are post-recombination and outflows from YSOs.

C. Formation in Very Dense Regions ( > 1012 cm3 s-1) - three-bodyreactions. e.g.,

3H H2 + H 2H + H2 2H2

with rate coefficients 10-32 – 10-33 cm6 s-1. Typical applicationsare cool stellar atmospheres, YSO disks and jets.

Formation of H22 (cont’d)

Photodissociation of H2As discussed in Lecture 18, H2 is dissociated in a two-step processwhen a far UV photon is absorbed in a Lyman or Werner transitionbelow 1108 Å, and the excited molecule decays into the continuum of the ground electronic state with the emission of a photon in the 1300-1700 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

yields N(H2) > 1014 cm-2. Thus absorption occurs mainly in the damping wings of the Lyman and Werner lines. The theory of H2photodissociation 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 191 375 1974) and Federman, Glassgold & Kwan(ApJ, 227, 466, 1979; FGH)

bN

cmef l

elu ππτ

2

)0( =

Photodissociation of H2 (cont’d)

[ ]∫∞

∞−

−−= ))(exp(1)( νϕννϕ jkjjk sNdW

FGH showed that, in the presence of H2 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:

Here σjk(∆ν)=sjk φ(ν) is the frequency dependent absorptioncross section. Differentiation leads to:

))()(

))(exp()(

))(exp()()(

νϕνϕ

νϕνϕ

νϕνϕ

jk

jkjjk

jkj

jk

sdv

sNsdv

NjsjkdvsN

W

∫

∫

∫

∞

∞−

∞

∞−

∞

∞−

−=

−=∂

∂

jk

jk

jkjjk

jkj

jk Jsdv

sNsdv

sNW

≡−

==∂

∂

∫

∫∞

∞−

∞

∞−

))()(

))(exp()(

)(νϕνϕ

νϕνϕ

This ratio is called the line self-shielding function Jjk

It gives the absorption cross section, averaged over the lineshape, taking into account the attenuation of the dissociatingradiation by H2 itself; it can be considered a function of the optical depth at line center

)broadeningDoppler for ()0()0(2

bN

cmefsN l

ejkjkj ππϕτ ==

Jjk is a decreasing function of optical depth (or column),starting from unity at zero optical depth. Having alreadyanalyzed the dependence of W on optical depth, FGKobtain the dependence of Jjk by differentiation.

Behavior of the Curve of Growth & Line Self-Shielding Function

There are three regimes of equivalent width, depending on optical depth at line center:

τ0 << 1 linear ~ τ0

τ0 > 1 flat or logarithmic ~ logτ01/2

τ0 >> 1 square-root ~ τ01/2

Differentiating W gives the corresponding behavior for J:τ0 << 1 J = 1τ0 > 1 J ~ 1/τ0τ0 >> 1 J ~ τ0

-1/2

In practice, the Lyman & Werner bands are usually optically thick, and the dissociation rate decreases as the inverse square root of the H2 column density

Follow FGK (or Warin et al. A&A 308 535 1996 for a more detailed theory). Balance the photodestruction rate G0Jn(H2) against the grain formation rate RnHn(H) while satisfying mass conservation, nH=n(H) + 2n(H2). The photo rate is calculated summing all accessible Lyman & Werner band transitions. The effects of other absorbers of 912-1108 Åradiation are also easy to include.

Abundance of H2 in Diffuse Clouds (and The Edges of Thick Clouds)

1110

2/11-31711

2H

02

s106,)K100/(scm103,cm 1075.6

)H(211)H(

2- −−− ×=×=×=

≅=+

=

GTRM

NMJ

RnJGx ε

ε

⎟⎟⎠

⎞⎜⎜⎝

⎛=

H

-3

H

100cmMyr101nRn

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

Spatial Variation of the H-H2 TransitionN(H2)

NH

dashed curve: Glassgold & Langer (ApJ 193 73 1973, similar to HWS) a. refined FGK theory (ground state)b. with dustc. thermal level populationd. collisional-radiative populationdotted-curve: 10% conversion

Limited observational support for the theory of H2 formation & destruction

dots - Copernicus dataline - FGK theory

4. Formation and Destruction of COA. Formation in Cool Regions: Fast ion-molecule reactions

k ~10-9 cm3 s-1

O+ + H2 → OH+ + H H3+ + O → OH+ + H2

OH+ + H2 → OH2+ + H or H3

+ + OH → OH2+ + H2

OH2+ + H2 → OH3

+ + H H3+ + H2O → OH3

+ + H2

and fast dissociative recombination, β ~10-9 cm3 s-1 :

OH+ + e → O + HOH2+ + e → OH + H and O + H2 or 2H 22% OHOH3+ + e → H2O + H and OH + H2 or 2H, etc. 75% OH

OH is a crucial intermediary in the formation of CO, and it requires H2 as a precursor. But the synthesis of OH is driven by the ionsO+ and H3

+ (to be discussed in Lecture 20). This scheme also produces H2O and other O-bearing molecules, as will also be discussed.

Gas-Phase Synthesis of COThe efficiency of ion-molecule synthesis of OH depends on the ratio ofreaction rates representing the competition between formation anddestruction, k x(H2)/ βx(e). In diffuse clouds with x(e) ~10-4 and x(H2) > 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 abundantCO+ + H2 → HCO+ + H than C+, the neutral reactionC+ + H2O → HCO+ + H is fast: C + OH → CO + HHCO+ + e → CO + H

The idea of using ion-molecule reactions to initiate low-temperatureinterstellar chemistry is due to WD Watson (ApJ 183 L17 1973) and Herbst & Klemperer (ApJ 185 505 1973), 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 188 255 1974).

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

117CR s105MeV) 2( −−×≈>ς

This refers to the ionization of atomic H, i.e, CR + H → H+ + e

Similar rates applies to CR ionization of He and H2,CR + He → He+ + e

CR + H2 → H2+ + e and H+ + H (5%)

But H2+ is very reactive:

H2+ + H → H+ + H2

H2+ + H2 → H3

+ + H2and leads to H3

+, the pivotal ion in the ion-molecule chemistry.

Destruction of CO1. In partially UV-shielded regions, the same two-step process operates as for H2, with the following changes:

a. G0 = 2x10-10 s-1 (for the mean UV radiation field),b. self-shielding is reduced since usually x(CO) << x(H2).c. mutual shielding in the 912-1108 Å band by H2 and even

some lines of atomic H are importnatd. mutual shielding of the isotopes 13C and C17O by the lines

of 12CO and H2 occurs2. In well-shielded (thick) clouds, CO is destroyed by the fast ion-

molecule reactionHe+ + CO → C+ + O + He.

B. Formation in Warm Regions - Radical reactions:

O + H2 --> OH + H (slow) C + OH --> CO + H (moderately fast)

Formation and Destruction of CO (cont’d)

Synthesis of CO: SummaryAll of the ion-molecule reactions for CO are fast, but theions may be in short supply in dense shielded regions because their abundances depend on ςCR/nH.

In partially shielded regions the time scales are short,but only a small fraction of carbon goes into CO becauseof strong photodissociation.

Slow formation of H2 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..