Download astrochemistry_caselli

School of Physics and Astronomy
FACULTY OF MATHEMATICS & PHYSICAL SCIENCES
Introduction to
QuickTime™ and a
decompressor
are needed to see this picture.
Astrochemistry
QuickTime™ and a
decompressor
are needed to see this picture.
Paola Caselli
Outline
• Astrochemical processes:
1.
2.
3.
4.
5.
6.
7.
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…
Interstellar Molecules
Known Interstellar Molecules (Total: 151 as of today)
Amino acetonitrile in SgrB2(N)
(Belloche et al. 2008)
H
H
O
C
C
N
H
H
O
H
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.
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):
RH 2
1
 n H v H A ng S H 
2
 1017 cm-3s -1
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
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.
Molecule Dissociation energy (eV)
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 ??
Once H2 is formed, the fun starts…
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 !
Dissociation
energy or
bond strength
Once H2 is formed, the fun starts…
Molecule Dissociation energy (eV)
H2
4.48
CH
OH
CH+
OH+
C + H2 
 CH + H
C+ + H2 
 CH+ + H
3.47
4.39
4.09
5.10
(endothermic by 1.01 eV)
(endothermic by 0.39 eV)
OH + H (endothermic by 0.09 eV)
O+ H2 
 OH+ + H
O+ + H2 
(exothermic by 0.62 eV!)
Some technical details:
Ion-Neutral reactions
A+ + B  C + +
D
Exothermic ion-molecule reactions do not possess activation
energy because of the strong long-range attractive force(Herbst
& Klemperer 1973; Anicich & Huntress 1986):
R
V(R) = -  e2/2R4
kLANGEVIN = 2 e(/)1/2
 10-9 cm3 s-1
independent on T
Some techincal details:
Neutral-Neutral reactions
Energy to break
the bond of the
reactant.
A + BC  AB + C
1 eV for endothermic reactions
E
0.1-1 eV for exothermic reactions
Energy released by
the formation of
the new bond.
kb T < 0.01 eV
in molecular clouds
Example: O + H2 
OH + H
(does not proceed in cold clouds)
Duley & Williams 1984,
Interstellar Chemistry;
Bettens et al. 1995, ApJ
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:
H
H
H2+ + H2  H3+ + H
H
The cosmic-ray ionization rate, 
  10-18 s-1 from the
known spectrum of high
energy cosmic rays.
 < 10-14 s-1 from thermal
equilibrium in diffuse
clouds.
 6x10-17 s-1 from thermal
equilibrium in dark clouds.
(Dalgarno 2006, PNAS)
(Tielens 2005, The Physics and Chemistry of the Interstellar Medium)
3. The chemistry initiated by H3+
H3 +
+
Once H3 is formed, a cascade of
reactions greatly enhance the
chemical complexity of the ISM.
O
OH+
H2
In fact, H3+ can easily donate a
proton and allow larger molecules
to build.
H2 O+
H2
H3 O+
e
Example 
OXYGEN CHEMISTRY (the
formation of water in the ISM)
H2 O
e e
OH
O
O2
O
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
H3 +
CH+
H2
CH2+
H2
e
CH2
CH3+
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.
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
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.
5. Nitrogen Chemistry
Nitrogen chemistry differs from that of oxygen and carbon:
N + H3+ 
 NH+ + H2
The N-chemistry
starts with a neutralneutral reaction (e.g.):
CH + N  CN + H
+
tN2 ~ 106 yr
N2 vs. CO
Chemistry in Photodissociation
Regions (PDRs)
Sternberg & Dalgarno 1995
CO
N2
Chemical Evolution?
Suzuki et al. 1992
Chemical Evolution?
Dust has to be taken into account!
Freeze-out vs. free-fall:
tdep


1
9
1

10
m
/T
(n

)
yr
X
H
2
nd ad v t
 3 1/ 2
7
1/ 2
t ff  
yr
  4 10 (n H )
32G 
Walmsley 1991
van Dishoeck et al. 1993
Evidences of freeze-out:
solid features
QuickTime™ and a
decompressor
are needed to see this picture.
Spitzer
from van Dishoeck et al. 2003
QuickTime™ and a
decompressor
are needed to see this picture.
Pontoppidan et al. 2007
Evidences of freeze-out:
the missing CO
C17O(1-0) emission
(Caselli et al. 1999)
CO
hole
dust
peak
0.05 ly
Dust emission in a pre-stellar core
(Ward-Thompson et al. 1999)
Molecules
freeze out
onto dust
grains in the
center of
pre-stellar
cores 
Dust grain
Evidences of freeze-out:
deuterium fractionation
N2D+(2-1)
N2H+(1-0)
Dust emission in the pre-stellar core
L1544 (Ward-Thompson et al. 1999)
D-fractionation increases
towards the core center
(~0.2; Caselli et al. 2002;
Crapsi et al. 2004, 2005)
Evidences
of freeze-out:
6. Deuterium
fractionation
deuterium fractionation
H2D+ / H3+ (and D/H) increases:
(i) in cold gas
CO/H2 
H3+ + HD  H2D+ + H2 + 230 K
(ii) when the abundance of gas
phase neutral species decreases.
(Dalgarno & Lepp 1984)
Roberts, Millar & Herbst 2003
H2D+ +
N2  N2D+ + H2
CO  DCO+ + H2
D/H  0.3 !
Evidences
of freeze-out:
H2D+ in L1544
deuterium fractionation
Vastel et al. 2006
Caselli et al. 2003, 2008
oH 2D +
CSO
N2H+(1-0)
IRAM
N2D+(2-1)
IRAM
Evidences
of freeze-out:
D-fractionation
and ion fraction
deuterium fractionation
ge-
PAH-, PAH
/n(H2)
neutrals
H3+
ege-
g-
HD
H2
H2D+
HD
H2
D2H+
HCO+, N2H+…
Wootten et al. 1979
Guelin et al. 1982
Bergin et al. 1998
Caselli et al. 1998
Dalgarno 2006
neutrals
DCO+, N2D+…
PAHPAH
neutrals
DCO+, N2D+…
Uncertainties:
1/3 * PAHs, PAH-s
* neutrals (O)
* ortho:para H2
2/3
PAHPAH
eg-
HD
H2
D3+
neutrals
PAHPAH
DCO+, N2D+…
QuickTime™ and a
decompressor
are needed to see this picture.
PAHs
What happens after a protostar is born?
e
are ne
and a
Time™
Quick ompress or icture.
dec ee this p
s
ded to
QuickTime™ and a
decompressor
are needed to see this picture.
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)
are n
ee
Quic
kT
de ime™
ded t compres and a
so
o see
this p r
icture
.
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)
SO
HCN
Quic
kTim
are n
e
eede decompr ™ and a
d to s
e
ee th ss or
is pic
ture.
HCO+
HCOOCH3
CH3OH
CH3CH2CN
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
are n
ee
Quic
kT
de ime™
ded t compres and a
so
o see
this p r
icture
.
What happens after a protostar is born?
QuickTime™ and a
decompressor
are needed to see this picture.
• dust heating, X-rays nearby protostars
(mantle processing and evaporation)
• dust (mantles and cores) sputtering +
vaporization along protostellar outflows
7. Surface Chemistry
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)
7. Surface Chemistry
REACTANTS: MAINLY MOBILE ATOMS AND RADICALS
A +
B 
AB
association
H +
H  H2
H +
X  XH (X = O, C, N, CO, etc.)
WHICH CONVERTS
Accretion
10/[Tk1/2 n(H2)] days
O  OH  H2O
C  CH  CH2  CH3  CH4
N  NH  NH2  NH3
CO  HCO  H2CO  H3CO  CH3OH
Diffusion+Reaction
tqt(H) 10-5-10-3 s
Watson & Salpeter 1972; Allen & Robinson 1977; Pickes & Williams 1977;
d’Hendecourt et al. 1985; Hasegawa et al. 1992; Caselli et al. 1993
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
QuickTime™ and a
decompressor
are needed to see this picture.
Henning & Semenov 2008
What happens in
protoplanetary disks?
Chemical Structure of PPDs
UV, X-rays
                         
surface
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?)
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
HCO+(3-2)
DCO+(3-2)
first image!
DCO+/HCO+
and
DCN/HCN
~0.05 in
TW Hydrae
HCN(3-2)
DCO+/HCO+~
0.05 in L1544
DCN(3-2)
first image!
ALMA
is needed !
Links to the Solar System ?
HDO IN THE DISK OF DM Tau:
HDO/H2O~0.01
SOURCE
Class 0 protostars (HC)
Protoplanetary disks
Comets
Carbonaceous
chondrites
Ground transition at 464 GHz Oceans
with JCMT
(Ceccarelli et al. 2005)
Herschel is needed !
HDO/H2O
~0.03
~0.01?
~3.0x10-4
~1.510-4
1.610-4
Links to the Solar System ?
The assemblage of planets…
Links to the Solar System ?
Chondrites: interstellar ovens?
cement
condrule
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 ?
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.)
L-Alanine
L-Aspartic Acid
L-Glutamine
Glycine
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
Spectroscopic
Chemical Analysis of
Atmophere.
Courtesy: Prof. G.W. Marcy, University of California, Berkeley
H2O vapor.
Methane:
Disequilibrium
chemicals.
CH4 + O2 --> CO2 +H2O
17
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.
QuickTime™ and a
decompressor
are needed to see this picture.
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 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
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
Constant need of interaction with real chemists (theory +
lab, gas-phase+solid state), who provide rate coefficients,
collisional rates, transition frequencies …