Introduction to Astrochemistry Paola Caselli School of Physics and Astronomy FACULTY OF MATHEMATICS & PHYSICAL SCIENCES.

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Introduction

to

Astrochemistr

y

Paola CaselliPaola Caselli





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?


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 !


(endothermic by 1.01 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,


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)


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



Evidences of freeze-out:solid features

from van Dishoeck et al. 2003

Pontoppidan et al. 2007



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





PAHsPAHs

What happens after a protostar is born?



Qu

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and a

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re.


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)




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+

HCOOCH3CH3OH CH3CH2CN

SO


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






• 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?



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


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.


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

TIFF (LZW) decompressorare needed to see this picture.

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


‘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!


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.)


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



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

Introduction to Astrochemistry Paola Caselli School of Physics and Astronomy FACULTY OF MATHEMATICS & PHYSICAL SCIENCES.

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