1. Introduction to Astrochemistry
School of Physics and Astronomy
FACULTY OF MATHEMATICS & PHYSICAL SCIENCES
Introduction to Astrochemistry
CH3COCH3
Interstellar molecules are crucial because:
they are observational tools,
they influence star formation,
they are the building blocks of life.
Paola Caselli

2. Outline Astrochemical processes:
The formation of H2
H3+ formation
The chemistry initiated by H3+
Formation and destruction of CO
Nitrogen chemistry
Deuterium fractionation
Surface chemistry
Examples: pre-stellar cores, protostellar envelopes, outflows, hot cores, protoplanetary disks…

3. Interstellar Molecules
Known Interstellar Molecules (Total: 151 as of today) C O N H
Glycine - the simplest amino acid
Amino acetonitrile in SgrB2(N)
(Belloche et al. 2008)

4. How do molecules form in the interstellar medium ?
The most elementary chemical reaction is the association of A and B to form a molecule AB with internal energy:
A + B  AB*
The molecule AB* must loose the internal energy. In the Earth atmosphere, where the number of particles per cubic centimeter (cc) is very large (~1019), the molecule looses its energy via three-body reactions:
Three body reactions do not play a significant role in gas-phase interstellar chemistry
AB* + M  AB
But this is not an efficient process in interstellar clouds (FAB~10-36n3cm-3s-1), where the number of particles per cc ranges between a few hundred and 107.

5. 1. The formation of H2
The reaction that starts the chemistry in the interstellar medium is the one between two hydrogen atoms to form molecular hydrogen:
H + H  H2
This reaction
happens on the
surface of dust
grains.

6. 1. The formation of H2
The H2 formation rate (cm-3 s-1) is given by (Gould & Salpeter 1963;
Hollenbach & Salpeter 1970; Jura 1974; Pirronello et al. 1999; Cazaux & Tielens 2002; Habart et al. 2003; Bergin et al. 2004; Cuppen & Herbst 2005):
H- + H -> H + e (R~10^{-21} cm^{-3} s^{-1})
Bergin et al. 2004: behind large scale flows, the timescale for H2 formation is ~10 Myr
nH gas number density
vH H atoms speed in gas-phase
A  grain cross sectional area
ng dust grain number density
SH sticking probability
  surface reaction probability

7. Once H2 is formed, the fun starts…
H2 is the key to the whole of interstellar chemistry. Some important species that might react with H2 are C, C+, O, N… To decide whether a certain reaction is chemically favored, we need to examine internal energy changes.
Dissociation energy (eV)
Molecule H2
4.48 CH
3.47 OH
4.39
CH+
4.09
OH+
5.10
Question: Can the following reactions proceed in the cold
interstellar medium?
C + H2  CH + H ??
C+ + H2  CH+ + H ??
O + H2  OH + H ??
O+ + H2  OH+ + H ??

8. Once H2 is formed, the fun starts…
Dissociation energy or bond strength
C + H2  CH + H ??
4.48 eV
3.47 eV
The bond strength of H2 is larger than that of CH the reaction is not energetically favorable.
The reaction is endothermic (by = 1.01 eV) and cannot proceed in cold clouds, where kb T < 0.01 eV !
If T = 100 -> k_bT~0.01 eV

9.  Once H2 is formed, the fun starts… O+ H2  OH + H 
Dissociation energy (eV)
Molecule H2
4.48 CH
3.47 OH
4.39
CH+
4.09
OH+
5.10
(endothermic by 1.01 eV)
(endothermic by 0.39 eV)
(endothermic by 0.09 eV)
C + H2  CH + H
C+ + H2  CH+ + H
O+ H2  OH + H 
If T = 100 -> k_bT~0.01 eV
O+ + H2  OH+ + H 
(exothermic by 0.62 eV!)

10. A+ + B  C+ + D Some technical details: Ion-Neutral reactions R
Exothermic ion-molecule reactions do not possess activation energy because of the strong long-range attractive force(Herbst & Klemperer 1973; Anicich & Huntress 1986): R
If T = 100 -> k_bT~0.01 eV
V(R) = -  e2/2R4
kLANGEVIN = 2 e(/)1/2
 10-9 cm3 s-1
independent on T

11.  A + BC  AB + C E  kb T < 0.01 eV Some techincal details:
Neutral-Neutral reactions
A + BC  AB + C
Energy to break the bond of the reactant.
1 eV for endothermic reactions
E 
0.1-1 eV for exothermic reactions
kb T < eV
in molecular clouds
Energy released by the formation of the new bond.
If T = 100 -> k_bT~0.01 eV
Example: O + H2  OH + H
(does not proceed in cold clouds) 
Duley & Williams 1984,
Interstellar Chemistry;
Bettens et al. 1995, ApJ

12. 2. Cosmic-ray ionization of H2
After the formation of molecular hydrogen, cosmic rays ionize H2 initiating fast routes towards the formation of complex molecules in dark clouds:
H2 + c.r.  H2+ + e- + c.r.
Once H2+ is formed (in small percentages), it very quickly reacts with the abundant H2 molecules to form H3+, the most important molecular ion in interstellar chemistry:
ION - MOLECULE REACTION H
H2+ + H2  H3+ + H

13. The cosmic-ray ionization rate, 
  s-1 from the
known spectrum of high
energy cosmic rays.
 < s-1 from thermal
equilibrium in diffuse
clouds.
 6x10-17 s-1 from thermal
equilibrium in dark clouds.
(Dalgarno 2006, PNAS)
PNAS = Proceedings of the National Academy of Science
(Tielens 2005, The Physics and Chemistry of the Interstellar Medium)

14. 3. The chemistry initiated by H3+
OH+
H3O+
H2O+ O H2
H2O OH e O2
Once H3+ is formed, a cascade of reactions greatly enhance the chemical complexity of the ISM.
In fact, H3+ can easily donate a proton and allow larger molecules to build.
Example 
OXYGEN CHEMISTRY (the formation of water in the ISM)

15. 3. The chemistry initiated by H3+
CARBON CHEMISTRY (the formation of hydrocarbons)
The formation of more complicated species from neutral atomic carbon begins with a sequence very similar to that which starts the oxygen chemistry: C
CH+
CH2+
CH3+
CH2
H3+ H2 e CH
A. Proton transfer from H3+ to a neutral atom;
B. Hydrogen abstraction reactions terminating in a molecular ion that does
not react with H2;
C. Dissociative recombination with electrons.

16. 4. Formation and destructio1n of CO
[a] C + H3O+  HCO+ + H2
[b] O + CH3+  HCO+ + H2
[c] HCO+ + e  CO + H is the most important source of CO.
CO is very stable and difficult to remove. It reacts with H3+:
[d] H3+ + CO  HCO+ + H2
but reaction [c] immediately reform CO.
The main mechanisms for removing CO are:
[e] He+ + CO  He + C+ + O
[f] h + CO  C + O
Some of C+ react with OH and H2O (but not with H2):
[g] C+ + OH  CO+ + H
[h] CO+ + H2  HCO+ + H
[i] C+ + H2O  HCO+ + H

17. The timescale to form CO
Assume: dark region where all H is in H2 and all atoms more
massive than He are in neutral atomic form.
The timescale on which almost all carbon becomes contained in CO (nO > nC) is at least equal to the timescale for one hydrogen molecule to be ionized for every C: nC/[ n(H2)] = 2 nC/[ nH]
For  = 610-17 s-1 and nC/nH = 10-4, the above expression gives a value of 105 yr.

18.  5. Nitrogen Chemistry CH + N  CN + H tN2 ~ 106 yr
Nitrogen chemistry differs from that of oxygen and carbon:
N + H3+  NH+ + H2
The N-chemistry
starts with a neutral-
neutral reaction (e.g.):
CH + N  CN + H  +
tN2 ~ 106 yr

19. N2 vs. CO Chemistry in Photodissociation Regions (PDRs)
Sternberg & Dalgarno 1995 CO N2
The Orion Bar

20. Chemical Evolution?
Suzuki et al. 1992

21. Dust has to be taken into account! Freeze-out vs. free-fall:
Chemical Evolution?
Dust has to be taken into account!
Freeze-out vs. free-fall:
Walmsley 1991
van Dishoeck et al. 1993

22. Evidences of freeze-out: solid features
Spitzer
from van Dishoeck et al. 2003
Pontoppidan et al. 2007

23. Evidences of freeze-out: the missing CO
C17O(1-0) emission
(Caselli et al. 1999) CO
hole
Dust grain
Molecules freeze out onto dust grains in the center of pre-stellar cores 
dust
peak
0.05 ly
Dust emission in a pre-stellar core
(Ward-Thompson et al. 1999)

24. Evidences of freeze-out: deuterium fractionation
N2D+(2-1)
N2H+(1-0)
D-fractionation increases towards the core center
(~0.2; Caselli et al. 2002; Crapsi et al. 2004, 2005)
Dust emission in the pre-stellar core
L1544 (Ward-Thompson et al. 1999)

25. Evidences of freeze-out: deuterium fractionation
H2D+ / H3+ (and D/H) increases:
(i) in cold gas
D/H  0.3 !
CO/H2 
H3+ + HD  H2D+ + H K
(ii) when the abundance of gas phase neutral species decreases.
(Dalgarno & Lepp 1984)
Roberts, Millar & Herbst 2003
N2  N2D+ + H2
H2D+ +
CO  DCO+ + H2

26. Evidences of freeze-out: deuterium fractionation H2D+ in L1544
Vastel et al. 2006
Caselli et al. 2003, 2008
o-H2D+
CSO
N2H+(1-0)
IRAM
N2D+(2-1)
IRAM

27. Evidences of freeze-out: deuterium fractionation
D-fractionation and ion fraction g- e-
PAH-, PAH
Wootten et al. 1979
Guelin et al. 1982
Bergin et al. 1998
Caselli et al. 1998
Dalgarno 2006
H3+
neutrals
/n(H2)
HCO+, N2H+… HD H2 e-
neutrals
Uncertainties:
* PAHs, PAH-s
* neutrals (O)
* ortho:para H2
DCO+, N2D+…
H2D+
1/3 g-
PAH-
PAH HD H2 e-
neutrals
DCO+, N2D+…
D2H+
2/3 g-
PAH-
PAH HD H2 e-
neutrals
DCO+, N2D+…
D3+ g-
PAH-
PAHs
PAH

28. What happens after a protostar is born?

29. D2CO/H2CO = 0.1 CHD2OH/CH3OH = 0.02 D2S/H2S = 0.02 ND3/NH3 = 0.001
What happens after a protostar is born?
Large abundances of multiply deuterated species in (Class 0) protostellar envelopes (Ceccarelli et al. 1998; Parise et al. 2002, 2004, 2006; van der Tak et al. 2002; Vastel et al. 2003)
D2CO/H2CO = 0.1
CHD2OH/CH3OH = 0.02
D2S/H2S = 0.02
ND3/NH3 = 0.001
CD3OH/CH3OH = 0.02

30. What happens after a protostar is born?
Complex organic molecules in hot cores and hot corinos
(e.g. Wright et al. 1996; Cazaux et al. 2003; Bottinelli et al. 2004,2008; Kuan et al. 2004)
HCN
HCO+
HCOOCH3
CH3OH
CH3CH2CN SO

31. What happens after a protostar is born?
Strong H2O, SiO, CH3OH, NH3, emission (e.g. Bachiller 1996) and complex molecules (C2H5OH, HCOOCH3: Arce et al. 2008) along outflows.
Jørgensen et al. 2004
Although gas phase chemistry can explain H2O, it is difficult to understand the observed abundance of the other species.
At T > 200 K: O + H2 -> OH + H and OH + H2 -> H2O + H
(e.g. Hartquist et al. 1980)

32. What happens after a protostar is born?
dust heating, X-rays nearby protostars (mantle processing and evaporation)
dust (mantles and cores) sputtering + vaporization along protostellar outflows

33. 7. Surface Chemistry 106 sites
thermal hopping
quantum tunneling
106 sites
Tielens & Hagen (1982); Tielens & Allamandola (1987); Hasegawa et al. (1992); Tielens 1993;
Cazaux & Tielens (2002); Cuppen & Herbst (2005); Cazaux et al. (2008); Garrod (2008)

34. 7. Surface Chemistry A + B  AB association Accretion
REACTANTS: MAINLY MOBILE ATOMS AND RADICALS
A B  AB association
H H  H2
H X  XH (X = O, C, N, CO, etc.)
Accretion
WHICH CONVERTS
O  OH  H2O
C  CH  CH2  CH3  CH4
N  NH  NH2  NH3
CO  HCO  H2CO  H3CO  CH3OH
10/[Tk1/2 n(H2)] days
Diffusion+Reaction
tqt(H)  s
Watson & Salpeter 1972; Allen & Robinson 1977; Pickes & Williams 1977;
d’Hendecourt et al. 1985; Hasegawa et al. 1992; Caselli et al. 1993

35. What happens in protoplanetary disks?
Aikawa & Herbst 1999; Markwick & Charnley 2003; Aikawa & Nomura 2006; Bergin et al. 2007; Dutrey et al. 2007; Meijerink, Poelman et al. 2008; Semenov et al. 2008
Henning & Semenov 2008

36. Chemical Structure of PPDs
What happens in
protoplanetary disks?
Chemical Structure of PPDs
surface
UV, X-rays
intermediate                         
midplane      
UV, c.r.
Surface layer : n~104-5cm-3, T>50K
Photochemistry (high CN/HCN)
Intermediate : n~106-7cm-3, 20<T<40K
Dense cloud chemistry (freeze-out, D-fractionation)
Midplane : n>107cm-3, T<20K
Freeze-out (are parent-cloud species preserved?)

37. ALMA is needed ! What happens in protoplanetary disks?
DCO+ (van Dishoeck et al. 2003; Guilloteau et al. 2006)
and H2D+ (Ceccarelli et al. 2004) detected.
QI, WILNER, AIKAWA, BLAKE, HOGERHEIJDE 2008
DCO+/HCO+ and DCN/HCN ~0.05 in
TW Hydrae
DCO+/HCO+~0.05 in L1544
HCO+(3-2)
DCO+(3-2)
first image!
HCN(3-2)
ALMA is needed !
DCN(3-2)
first image!

38. Links to the Solar System ?
HDO IN THE DISK OF DM Tau: HDO/H2O~0.01
SOURCE
HDO/H2O
Class 0 protostars (HC)
~0.03
Protoplanetary disks
~0.01?
Comets
~3.0x10-4
Carbonaceous
chondrites
~1.510-4
Oceans
1.610-4
Ground transition at 464 GHz with JCMT
(Ceccarelli et al. 2005)
Herschel is needed !

39. Links to the Solar System ?
The assemblage of planets…

40. Links to the Solar System ?
Chondrites: interstellar ovens?
cement
Most are from undifferentiated parent bodies: original build-up particles still recognizable:
Chondrules (mm size spherules)
Matrix (`cement’ between chondrules: <10 m particles)
Calcium-Aluminium-rich Inclusions (CAIs, cm size, rare)
condrule

41. Links to the Solar System ?
‘Cement’ between chondrules:
Consists of tiny particles (~ interstellar dust)
Often contains water and carbon
Often contains hydrous minerals resulting from ancient interaction of liquid water and primary minerals.
Must have been liquid water in planetesimals!

42. Links to the Solar System ?
Carbonaceous chondrites contain a substantial amount of C, up to 3% by weight.
~70 amino acids have been identified in carbonaceous chondrites; 8 of these are found in terrestrial proteins
(Botta & Bada 2002, Survey Geophys.)
Point out that we do not know how these molecules are formed. Still unexplored chemistry!
Most amino acids can exist in either of two optical isomers. While L-amino acids represent the vast majority of amino acids found in proteins, D-amino acids are found in some proteins produced by exotic sea-dwelling organisms, such as cone snails.[21]
The left hand is a non-superimposable mirror image of the right hand; no matter how the two hands are oriented, it is impossible for all the major features of both hands to coincide.
L-Alanine
L-Aspartic Acid
L-Glutamine
Glycine

43. Exoplanets
Brown Dwarf 2M1207 and its planetary companion (~14 MJ; Chauvin et al. 2005).

44. Exoplanets Initial studies of hot Jupiters’ atmospheres:
Richardson et al (Spitzer)
de Mooij et al (WHT+UKIRT)
Sing & Lopez-Morales 2009 (Magellan+VLT)
WHT = William Herschel Telescope

45. Exoplanets Darwin & TPF will detect Biomarkers: O3 H2O vapor.
Spectroscopic
Chemical Analysis of Atmophere.
Methane:
Disequilibrium chemicals.
CH4 + O2 --> CO2 +H2O
Courtesy: Prof. G.W. Marcy, University of California, Berkeley 17

46. Exoplanets
NASA's Kepler spacecraft, scheduled to launch in March on a journey to search for other Earths, has arrived in Cape Canaveral, FL
For four years, Kepler will monitor 100,000 stars in our Galaxy,looking for (Earthlike) planetary transits.

47. Summary Prestellar cores: CN, N2H+, NH3, N2D+, DCO+, o-H2D+…
Ion-molecule reactions, freeze-out, deuterium fractionation,
surface chemistry
Outflows: H2O, CH3OH, NH3, SiO, S-bearing species
Grain sputtering, grain-grain collisions, neutral-neutral reactions
Hot Cores: CH3CN, HCOOCH3, complex saturated molecules
Grain mantle evaporation, neutral-neutral reactions, surface chemistry
PP Disks: CO, CN, HCN,N2H+,HCO+, DCO+, o-H2D+
Ion-molecule reactions, freeze-out, D-fractionation, surface
chemistry, photochemistry, X-rays, dust coagulation

48. Main Uncertainties Cosmic-ray ionization rate
Elemental abundance in dark clouds (e.g. metals)
Oxygen chemistry ( Herschel)
PAHs abundance
Surface chemistry and gas phase high-T chemistry
H2 ortho-to-para ratio
Say something about the problems with numerical treatment of surface chemistry
Constant need of interaction with real chemists (theory +
lab, gas-phase+solid state), who provide rate coefficients,
collisional rates, transition frequencies …