Introduction to Astrochemistry Paola Caselli School of Physics and Astronomy FACULTY OF MATHEMATICS...
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Transcript of Introduction to Astrochemistry Paola Caselli School of Physics and Astronomy FACULTY OF MATHEMATICS...
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Introduction
to
Astrochemistr
y
Paola CaselliPaola Caselli
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School of Physics and AstronomyFACULTY OF MATHEMATICS & PHYSICAL SCIENCES
Outline
• Astrochemical processes:1. The formation of H2
2. H3+ formation
3. The chemistry initiated by H3+
4. Formation and destruction of CO5. Nitrogen chemistry6. Deuterium fractionation7. Surface chemistry
1.Examples: pre-stellar cores, protostellar envelopes, outflows, hot cores, protoplanetary disks…
Interstellar MoleculesKnown Interstellar Molecules (Total: 151 as of today)
Amino acetonitrile in SgrB2(N) (Belloche et al. 2008)
C
C
O
O
N
H
HH
HH
Glycine - the simplest amino acid
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:
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.
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.
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):
€
RH2=
1
2nHvHAngSHγ
≅ 10−17 cm-3s-1
nH gas number densityvH H atoms speed in gas-phaseA grain cross sectional areang dust grain number densitySH sticking probability surface reaction probability
1. The formation of H2
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.
H2 4.48
CH 3.47
OH 4.39
CH+ 4.09
OH+ 5.10
Dissociation energy (eV)Molecule
C + H2 CH + H ??C+ + H2 CH+ + H ??
O + H2 OH + H ??O+ + H2 OH+ + H ??
Question: Can the following reactions proceed in the cold interstellar medium?
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 4.48-3.47 = 1.01 eV) and cannot proceed in cold clouds, where kb T < 0.01 eV !
Once H2 is formed, the fun starts…
(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
O+ + H2 OH+ + H (exothermic by 0.62 eV!)
H2 4.48
CH 3.47
OH 4.39
CH+ 4.09
OH+ 5.10
Dissociation energy (eV)Molecule
Some technical details: Ion-Neutral reactions
AA++ + B + B C C++ + D + D
Exothermic ion-molecule reactions do not possess activation energy because of the strong long-range attractive force(Herbst & Klemperer 1973; Anicich & Huntress 1986):
V(R) = - e2/2R4
R
kLANGEVIN = 2 e(/)1/2
10-9 cm3 s-1
independent on T
A + BC A + BC AB + C AB + C 1 eV for endothermic reactionsE 0.1-1 eV for exothermic reactions
kb T < 0.01 eV
in molecular clouds
Energy to break the bond of the reactant.
Energy released by the formation of the new bond.
Example: O + H2 OH + H(does not proceed in cold clouds)
Duley & Williams 1984, Interstellar Chemistry; Bettens et al. 1995, ApJ
Some techincal details: Neutral-Neutral reactions
2. Cosmic-ray ionization of H2
H2 + c.r. H2+ + e- + c.r.
After the formation of molecular hydrogen, cosmic rays ionize H2 initiating fast routes towards the formation of complex molecules in dark clouds:
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:
H2+ + H2 H3
+ + H H H
H
10-18 s-1 from the known spectrum of highenergy cosmic rays.
< 10-14 s-1 from thermalequilibrium in diffuse clouds.
6x10-17 s-1 from thermalequilibrium in dark clouds.
(Dalgarno 2006, PNAS)
(Tielens 2005, The Physics and Chemistry of the Interstellar Medium)
The cosmic-ray ionization rate,
3. The chemistry initiated by H3+
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.
H3+
OH+
H3O+
H2O+
O
H2
H2
H2O OH O
eee
O2
O
Example OXYGEN CHEMISTRY (the formation of water in the ISM)
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:
CH
C CH+ CH2+ CH3+
CH2H3+ H2 H2
e
e
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.
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
The timescale to form CO
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 = 610-17 s-1 and nC/nH = 10-4, the above expression gives a value of 105 yr.
Assume: dark region where all H is in H2 and all atoms more massive than He are in neutral atomic form.
5. Nitrogen Chemistry
Nitrogen chemistry differs from that of oxygen and carbon:
N + H3+ NH+ + H2
The N-chemistrystarts with a neutral-neutral reaction (e.g.):
CH + N CN + H
+
tN2 ~ 106 yr
N2 vs. CO
Chemistry in Photodissociation Regions (PDRs)Sternberg & Dalgarno 1995
CO
N2
The Orion Bar
Chemical Evolution?
Suzuki et al. 1992
Chemical Evolution?
Dust has to be taken into account!Freeze-out vs. free-fall:
Walmsley 1991van Dishoeck et al. 1993
€
tdep =1
αndπad2v t
≈109 mX /T (nHα )−1 yr
€
t ff =3π
32Gρ
⎛
⎝ ⎜
⎞
⎠ ⎟
−1/ 2
= 4 ×107 (nH )−1/ 2 yr
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Evidences of freeze-out:solid features
from van Dishoeck et al. 2003
Pontoppidan et al. 2007
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Spitzer
Evidences of freeze-out:the missing CO
C17O(1-0) emission(Caselli et al. 1999)
CO hole
dust peak
Dust emission in a pre-stellar core(Ward-Thompson et al. 1999)
0.05 ly
Molecules freeze out onto dust grains in the center of pre-stellar cores
Dust grain
Evidences of freeze-out: deuterium fractionation
D-fractionation increases towards the core center (~0.2; Caselli et al. 2002; Crapsi et al. 2004, 2005)
N2D+(2-1)
N2H+(1-0)
Dust emission in the pre-stellar core L1544 (Ward-Thompson et al. 1999)
Evidences of freeze-out: deuterium fractionation
H3+ + HD H2D+ + H2 + 230 K
6. Deuterium fractionation
D/H 0.3 !
CO/H2
(ii) when the abundance of gas phase neutral species decreases. (Dalgarno & Lepp 1984) Roberts, Millar & Herbst 2003
N2 N2D+ + H2
H2D+ +
CO DCO+ + H2
H2D+ / H3+ (and D/H) increases:
(i) in cold gas
Evidences of freeze-out: deuterium fractionation H2D+ in L1544
o-H2D+
CSO
N2D+(2-1)IRAM
N2H+(1-0)IRAM
Vastel et al. 2006
Caselli et al. 2003, 2008
Evidences of freeze-out: deuterium fractionation D-fractionation and ion fraction
Wootten et al. 1979Guelin et al. 1982Bergin et al. 1998Caselli et al. 1998
Dalgarno 2006
H3+
H2D+
/n(H2)neutrals
neutrals
e-g-
PAH-, PAH
HD H2e-
PAH-
D2H+
HD H2
HCO+, N2H+…
DCO+, N2D+…
D3+
HD H2
g-
e-
g-
e-
g-
PAH
neutrals
PAH-
DCO+, N2D+…
PAH
neutrals
PAH-
DCO+, N2D+…
PAH
1/3
2/3
Uncertainties:Uncertainties:* PAHs, PAH* PAHs, PAH--ss* neutrals (O)* neutrals (O)* ortho:para H* ortho:para H22
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PAHsPAHs
What happens after a protostar is born?
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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)
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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)
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HCN
HCO+
HCOOCH3CH3OH CH3CH2CN
SO
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
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What happens after a protostar is born?
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• dust heating, X-rays nearby protostars (mantle processing and evaporation)
• dust (mantles and cores) sputtering + vaporization along protostellar outflows
106 sitesTielens & Hagen (1982); Tielens & Allamandola (1987); Hasegawa et al. (1992); Tielens 1993;Cazaux & Tielens (2002); Cuppen & Herbst (2005); Cazaux et al. (2008); Garrod (2008)
7. Surface Chemistry
quantum tunneling
thermal hopping
7. Surface Chemistry REACTANTS: MAINLY MOBILE ATOMS AND RADICALS
A + B AB associationH + H H2
H + X XH (X = O, C, N, CO, etc.)
WHICH CONVERTS O OH H2O
C CH CH2 CH3 CH4
N NH NH2 NH3
CO HCO H2CO H3CO CH3OH
Watson & Salpeter 1972; Allen & Robinson 1977; Pickes & Williams 1977; d’Hendecourt et al. 1985; Hasegawa et al. 1992; Caselli et al. 1993
Accretion
Diffusion+Reaction
10/[Tk1/2 n(H2)] days
tqt(H) 10-5-10-3 s
What happens in protoplanetary disks?
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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
What happens in protoplanetary disks?
Chemical Structure of PPDs
Surface layer : n~104-5cm-3, T>50KPhotochemistry (high CN/HCN)
Intermediate : n~106-7cm-3, 20<T<40KDense cloud chemistry (freeze-out, D-fractionation)
Midplane : n>107cm-3, T<20KFreeze-out (are parent-cloud species preserved?are parent-cloud species preserved?)
UV, X-rayssurface
intermediate
midplane
UV, c.r.
What happens in protoplanetary disks?
DCO+ (van Dishoeck et al. 2003; Guilloteau et al. 2006) and H2D+ (Ceccarelli et al. 2004) detected.
HCO+(3-2)
DCO+(3-2) first image!
DCO+/HCO+
and DCN/HCN ~0.05 in TW Hydrae
DCO+/HCO+~0.05 in L1544
HCN(3-2)
DCN(3-2)first image!
QI, WILNER, AIKAWA, BLAKE, HOGERHEIJDE 2008
ALMAis needed !
Links to the Solar System ? HDO IN THE DISK OF DM Tau:
HDO/H2O~0.01QuickTime™ and a
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Ground transition at 464 GHz with JCMT
(Ceccarelli et al. 2005)
SOURCE HDO/H2O
Class 0 protostars (HC) ~0.03
Protoplanetary disks ~0.01?
Comets ~3.0x10-4
Carbonaceous
chondrites
~1.510-4
Oceans 1.610-4
Herschel is needed !
The assemblage of planets…
Links to the Solar System ?
Chondrites: interstellar ovens?
condrule
cement
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!
Links to the Solar System ?
L-Alanine L-Aspartic Acid L-Glutamine Glycine
~70 amino acids have been identified in carbonaceous chondrites; 8 of these are found in terrestrial proteins (Botta & Bada 2002, Survey Geophys.)
Links to the Solar System ?
Carbonaceous chondrites contain a substantial amount of C, up to 3% by weight.
Exoplanets
Brown Dwarf 2M1207 and its planetary companion (~14 MJ; Chauvin et al. 2005).
Exoplanets Initial studies of hot Jupiters’ atmospheres:
Richardson et al. 2007 (Spitzer)de Mooij et al. 2009 (WHT+UKIRT)Sing & Lopez-Morales 2009 (Magellan+VLT)
Exoplanets Darwin & TPF will detect Biomarkers:
O3 H2O vapor.
SpectroscopicChemical Analysis of Atmophere.
Methane:Disequilibrium chemicals.CH4 + O2 --> CO2 +H2O
17Courtesy: Prof. G.W. Marcy, University of California, Berkeley
Exoplanets
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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.
http://planetquest.jpl.nasa.gov/news/keplerArrival.cfm
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 speciesGrain sputtering, grain-grain collisions, neutral-neutral reactions
Hot Cores: CH3CN, HCOOCH3, complex saturated moleculesGrain 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
Constant need of interaction with real chemists (theory + lab, gas-phase+solid state), who provide rate coefficients,collisional rates, transition frequencies …
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