Interstellar chemistry

Liv Hornekær

Overall structure of the ISM

Overview picture

21 cm

2.6 mm

21 cm observations

Leiden-Dwingeloo/Argentina/Bonn

H atom

J=L+S F=I+J

La: 121.6 nm, 10.2 eV

From spectroscopy of the solar photosphere.

Relative abundance of the elements: A(X) = 106 (N

X/N

H)

2.6 mm

CO

J=0->1

CO

Overview of interstellar molecules

Molecules with Two Atoms

AlF AlCl C2 CH CH+ CN CO CO+ CP CS CSi HCl H2 KCl

NH NO NS NaCl OH PN SO SO+ SiN SiO SiS HF SH FeO N2

Molecules with Three Atoms

C3 C2H C2O C2S CH2 HCN HCO HCO+ HCS+ HOC+

H2O H2S HNC HNO MgCN MgNC N2H+ N20 NaCN OCS SO2 SiC2 CO2 NH2

H3+ AlNC

Molecules with Four Atoms

C3H C3H C3N C3O C3S C2H2 CH2D+ HCCN HCNH+ HNCO HNCS HOCO+

H2CO H2CN H2CS H3O+ SiC3 NH3

Molecules with Five Atoms

C5 C4H C4Si l-C3H2 c-C3H2 CH2CN CH4 HC3N HC2NC HCOOH H2CHN

H2C20 H2NCN HNC3 SiH4 H2COH+

Molecules with Six Atoms

C5H C5O C2H4 CH3CN CH3NC CH3OH CH3SH HC3NH+ HC2CHO HCONH2

H2C4 C5N

Molecules with Seven Atoms

C6H CH2CHCN CH3C2H HC5N HCOCH3 NH2CH3 C2H4O CH2CHOH

Molecules with Eight Atoms

CH3C3N HCOOCH3 CH3COOH C7H H2C6 CH2OHCHO CH2CHCHO

Molecules with Nine Atoms

CH3C4H CH3CH2CN (CH3)20 CH3CH20H HC7N C8H

Molecules with Ten Atoms

CH3C5N (CH3)2CO NH2CH2COOH CH3CH2CHO

Molecules with Eleven Atoms

HC9N

Molecules with Thirteen Atoms

HC11N

Chemical reactions in general

Reaction rate: k=Ae-Ea/kT

Stabilizing reaction products

Stabilizing reaction products

Radiation stabilization

Collisional stabilization:

AB*+ M AB + M

Chemical reactions in general

Reaction rate: k=Ae-Ea/kT

Reactions including radicals: Ea small

Reactions including ions: Ea ~ 0

Ion-molecule reactions

A+ + B -> C

+ + D

Typical rate: k=10-9

cm3/s

Rate equation:

= -k(T)n(A)n(B)

dn(A)

dt

Chemical models

~1000 gas phase reactions

~1-2 chemical reactions on dust grain surfaces.

A few models include ~40 surface reactions.

Reaction rates from:

1) Experimental measurements

2) Extrapolation of experimental measurements at high

temperature to interstellar temperatures

3) Theoretical calculations

4) Guess – or trial and error (Fitting the model to observations)

Experimental measurements

Ian Sims

CN + C2H6

Reaction rate: k=Ae-Ea/kT

Experimental measurements

Ian Sims

CN + C2H6

Reaction rate: k=Ae-Ea/kT

x

Carbon chemistry

• I.P. of C: 11.26 eV < 13.6 eV => majority of carbon as

C+

• C+ + H

2 → CH

2

+ + hν possible at low T (starting

reaction)

• CH2

+ => fast ion-molecule chemistry giving: CH, C

2,

…

• C+ + H

2 → CH

+ + H: endothermic by 0.4 eV

Average Interstellar Radiation

field

Sum of CMB (radio/FIR), thermal emission from dust (IR), cool stars + OB stars

(VIS, UV), hot ionized medium (FUV and X-ray)

Average Interstellar Radiation

field

Carbon chemistry

• I.P. of C: 11.26 eV < 13.6 eV => majority of carbon as

C+

• C+ + H

2 → CH

2

+ + hν possible at low T (starting

reaction)

• CH2

+ => fast ion-molecule chemistry giving: CH, C

2,

…

• C+ + H

2 → CH

+ + H: endothermic by 0.4 eV

Oxygen chemistry

I.P. of O: 13.618 > 13.598 eV => majority of oxygen as O

• Ionization by cosmic rays

• H2 or H + C.R. → H

2

+ or H

+ + C.R. + e

• H2

+ + H

2 → H

3

+ + H (fast)

• H+ or H

3

+ reacts with oxygen:

• H+ + O ↔ H + O

+ , O

+ + H

2 → OH

+ + H

• H3

+ + O → OH

+ + H

2

• When OH+ is formed fast ion-molecule reactions give OH, H

2O and

CO

• OH abundace is proportional to the ionization rate by cosmic

radiation: ζCR

=> observed OH abundance used to determine ζCR

Overview of interstellar molecules

Molecules with Two Atoms

AlF AlCl C2 CH CH+ CN CO CO+ CP CS CSi HCl H2 KCl

NH NO NS NaCl OH PN SO SO+ SiN SiO SiS HF SH FeO N2

Molecules with Three Atoms

C3 C2H C2O C2S CH2 HCN HCO HCO+ HCS+ HOC+

H2O H2S HNC HNO MgCN MgNC N2H+ N20 NaCN OCS SO2 SiC2 CO2 NH2

H3+ AlNC

Molecules with Four Atoms

C3H C3H C3N C3O C3S C2H2 CH2D+ HCCN HCNH+ HNCO HNCS HOCO+

H2CO H2CN H2CS H3O+ SiC3 NH3

Molecules with Five Atoms

C5 C4H C4Si l-C3H2 c-C3H2 CH2CN CH4 HC3N HC2NC HCOOH H2CHN

H2C20 H2NCN HNC3 SiH4 H2COH+

Molecules with Six Atoms

C5H C5O C2H4 CH3CN CH3NC CH3OH CH3SH HC3NH+ HC2CHO HCONH2

H2C4 C5N

Molecules with Seven Atoms

C6H CH2CHCN CH3C2H HC5N HCOCH3 NH2CH3 C2H4O CH2CHOH

H2 formation

No dipole allowed transitions

No radiation stabilization

H

3-body collisions – ok at high density

Diffuse/Dense cloud ISM densities => ~No 3-body collisions

Coalsack nebula in the

Southern Cross

H2 absorption bands in diffuse

clouds

HD110432 behind the Coalsack Nebula

ISM and star formation

H2

H2

H2 is an important cooling agent and

key to developing chemical complexity in ISM

- But how is H2 formed under ISM conditions?

202 nm

202 nm

Destruction mechanisms for H2

Lyman transition

Werner

transition

+ Cosmic

radiation

H2 formation rate from

observations

Typical diffuse cloud rate: k~3 10-17

n(H) cm3s-1

H2 formation

No dipole allowed transitions

No radiation stabilization

H

3-body collisions – ok at high density

Diffuse/Dense cloud ISM densities => ~No 3-body collisions

H2 formation in the gas phase

Radiative association (slow):

H + e- H

- + hn

Associative detachment:

H + H- H

2 + e

-

Typical diffuse cloud rate: k~10-21

n(H) cm3s-1

H2 formation rate: k~3 10

-17 n(H) cm

3s-1

Surface reactions

H

Surface reactions

Surface reactions

Objects: Dark nebulae

Molecules and dust grains

Depletion

D(X) = log(NX/N

H)-log(N

X/N

H)ISM

Relative depletion: d(X) = 1-10D(X)

=1-(NX/N

H) / (N

X/N

H)ISM

Measured in UV => diffuse and intercloud medium.

Absorption spectrum

H2O

H20

CH

CH

CO

Silicates

(Mg2SiO4 ,

Fe2SiO4)

Silicates

MgSiO3 (enstatite): Si-O stretch 9.7 mm O-Si-O bend 19.0 mm

Mg2SiO4 (fosterite): Si-O stretch 10.0 mm O-Si-O bend 19.5 mm

FeSiO3 (ferrosilite): Si-O stretch 9.5 mm O-Si-O bend 20.0 mm

SiC (Silicon carbide): Si-C stretch 11.2 mm

Carbon

Hydrogenated amorphous carbon: C-H stretch: 3.4 mm

Observed in the diffuse ISM

Measured towards Sgr A.

C60

+

C60 and C70 detected in

protoplanetary nebula Tc 1

C60

C70

Cami et al., Science 329, 1180, Sept. 2010

A few % of C

Diamond

ISO spectra of two pre-main-sequence stars

Lower curves are laboratory absorption

spectra for diamond nano-crystals

Aromatic hydrocarbon related

features

6.8 -10.8 microns

Aromatic hydrocarbon related

features

3.3 mm: C-H stretch

6.2 mm: C-C stretch

7.7 mm: C-C stretch

8.7 mm: C-H in plane bend

11.2 mm: C-H out of plane bend

6.8 -10.8 microns

PAH’er

Polycykliske Aromatiske Kulbrinter

Grænsen mellem molekyle og nano-partikel

Benzene Pyrene

Emission fra PAH’er

teori – 850 K PAH’er

Gennemsnitlig interstellar emission, => kun ~5% af C i PAH’er

PAH’er

Carbon

Interstellar carbon

Ice

Dust enshrouded protostar

H2O: O-H stretch: 3.05 mm

CO: C-O stretch: 4.67 mm

CH3OH: O-H stretch: 3.05 mm

H-O-H bend: 6.0 mm

C-H stretch: 3.53 mm

Water ice - morphology

Spectral lines – vibrational bands modified by local environment

Hindered rotation

Rotationel splitting

CO ice

Mixed ice

High shielding, low temperature: Even very volatile

molecules (e.g. CO) condenses out on dust grains.

Observations of 4.67 mm absorption in C-O

stretch, shows that CO and water are not mixed

Different molecules condenses at different temperatures

1. Large grains:

~100 nm

2. Surface structure or

smaller grains ?

4. PAHs or other

very small grains (VSG) < 10nm

3. Small carbon grains

< 20 nm

1 2 3

4

Greenberg 1996

Average interstellar extinction

curve

Dust grain sizes

Large grains: 20 nm - 1 mm

Silicate or carbon

Separate populations ? / composite grains ? /

Grains with carbon mantles ?

Very small grains: 1-20 nm,

Carbon

PAHs

Pre-solar dust grains in meteorites

Onion-like graphite particles

Silicates: Olivines (Mg2SiO

4, Fe

2SiO

4)

Atoms and molecules on surfaces

1

2a

2b Sticking

Sticking + hot atom

Scattering

Sticking:

S = Probability of 2a+2b

Sticking => adsorbed

Possible mechanisms for H2 formation

on surfaces

Desorption

Langmuir-Hinshelwood:

Adsorption

Diffusion

Recombination

Desorption

Eley-Rideal:

H(ads) + H(gas) H2(gas)

Hot Atom:

Strong/Weak

Adsorption Diffusion

Recombination

The effect on the ISM of H2 formation

v

Kinetic energy ?

n = 0

J = 0

Molecular excitation ?

Ereleased

~ 4.5 eV (50.000 oC)

=>

Dust grain heating ?

The effect on the ISM of H2

formation

Flores

Dashed line: No energy branching into kinetic energy

Curve 1: 0.5 eV in kinetic energy

Curve 2: 1.5 eV in kinetic energy

Curve 3: 2.25 eV in kinetic energy

Bringing the Interstellar medium to

a laboratory near you

Re-creating interstellare conditions ?

Interstellar pressure:

P= 10-13

atm

Interstellar temperatures:

T = 6-1000 K

Relevant surfaces:

Ice, graphite, PAHs, amorphous carbon, silicates

Surfaces of interstellar relevance

Water ice Atomic deposition

shield Cu

Graphite

T

H2 formation at high temperatures

– bare grains

Orion nebula – Orion bar

H on graphite

Eva Rauls

Neumann et al. Appl. Phys. A 55, 489 (1992)

Jeloica & Sidis, Chem. Phys. Lett. 300, 157 (1999)

Sha et al, Surface Science 496, 318 (2002)

H-Dimers on graphite

Dimer A

Dimer B

103 x 114 Å2

Vt = 884 mV, I

t = 0.16 nA

Dimers: Theori vs. Experiment

Vt = 884 mV, I

t = 0.16 nA

Vt=0.9 V, LDOS=1x10

-6 (eV)

-1 Å

-3

e. f.

Ortho dimer - Dimer A Para dimer - Dimer B

Pair formation

Hornekær et al. Phys. Rev. Lett. 97, 186102 (2006)

n=1 => First order

desorption

490 K => 1.4 eV

580 K => 1.6 eV

dQ

dt = -k0 e

- / T Qn kB EB

Zecho et al, J. Chem. Phys. 117, 8486 (2002)

H2 formation on graphite - TDS

H2 formation

Hornekær et al. Phys. Rev. Lett. 96, 156104 (2006)

Ortho Meta Para

Eley Rideal - Abstraction

Sha et al. (2002)

Zecho et al. (2002)

Matinazzo & Tantardini (2006)

Morisset et al. (2004)

Jeloaica & Sidis (2001)

Meijer et al. (2001)

Bachellerie et al. (2007)

Thomas et al. (2008)

H2 formation at lower temperatures

- ice covered surfaces

Water ice - morphology

Thermal Desorption Spectroscopy

HD from amorphous

water ice

D2 from graphite

Thermal Desorption Spectroscopy

HD from amorphous

water ice

D2 from graphite

Kinetic energy of formed H2

Laser desorption –

Time-of-flight

H D

QMS

ASW

Laser

Kinetic energy of D2 formed on

graphite

0 1.0 2.0 3.0 4.00

0.002

0.004

0.006

0.008

0.01

0.012

0.014

0.016

Translational Energy (eV)

D(E

)

~1.3 eV i translation

S. Baouche et al, J. Chem. Phys. 2006

Kinetic energy distribution

Baouche et al., J. Chem. Phys. 125, 084712 (2006 )

Kinetic energy of HD formed

on porous water ice

HD

D2

Solid line:

45 K Maxwell-Boltzmann velocity dist.

L. Hornekær et al., Science, 302, 1943 (2003)

H2 formed on porous water ice

Surface structure

determines energy

partitioning

Porous surface:

Grain heating

Slow H2

Result

Surface structure determines energy partitioning

Energy distribution in H2 formation

on graphite

L1630 Orion bar

Latimer et al., CPL 455, 174 (2008)

Dust grain morphology

Bare grains – porous and non-porous

Ice covered grains – maybe compacted by

formation/processing

Orion nebula – Orion bar

Energy distribution in H2 formation

and PDR observations

Observation show overpopulation in v=4

Does not fit shock and UV fluorescence models

- WHY?

Photo Dissociation Regions:

gas temperatures: 100 – 1000K

Blue: PAH emission

Green: H2 vibrational

line emission (FAST)

Red: CO emission

(30 m Telescope,

IRAM)

Star at NW

L1630 Orion bar

Energy distribution in H2 formation

and PDR observations

Observation show overpopulation in v=4

Does not fit Shock and UV fluorescense models

- formation pumping?

Photo Dissociation Regions:

gas temperatures: 100 – 1000K

Blue: PAH emission

Green: H2 vibrational

line emission (FAST)

Red: CO emission

(30 m Telescope,

IRAM)

Star at NW

L1630 Orion bar

H2 formation in different ISM

environments?

Hgas

+Hgas

H2

H + H + H

2 Problematic under

diffuse cloud conditions

Fine under

Dense cloud and PDR conditions

Hornekær et al., Science (2003)

Hornekær et al., PRL (2006)

H + H + H

2 Alternative candidate

Rauls and Hornekær (2007)

Hornekær et al., PRL (2006)

Katz et al., ApJ (1999)

Cuppen et al., MNRAS (2005)

PAHs as an H2 catalyst

• Correlations between high H2

formation rates and PAH

emission observed in PDRs with

low UV flux

– Habart et al.*

• PAH cations considered for H2

formation

– Snow et al.^

– LePage et al.#

*Habart et al., A&A, 397, 623 (2003)

Habart et al. A&A, 414, 531 (2004)

^Snow et al. Nature 391, 259 (1998)

#LePage et al. Ap.J., 704, 274, (2009)

Here:

the role of neutral PAHs?

H-PAH interaction

Density Functional Theory (DFT) calculations

reveal low barrier routes to

H-PAH formation and PAH catalyzed H2 formation

Rauls and Hornekær, Astrophys. J. 679, 531 (2008)

60 meV

-1.4 eV

C24H13 C24H12 + H

H-PAH interaction

Rauls and Hornekær, Astrophys. J. 679, 531 (2008)

Density Functional Theory (DFT) calculations

reveal low barrier routes to

H-PAH formation and PAH catalyzed H2 formation

H-PAH formation

Superhydrogenated PAHs

• Evidence in IR emission

– C-H stretching mode

• 3.3 μm – aromatic

• 3.4 μm – aliphatic

• High UV flux (Orion bar)

– Limited excess hydrogen

• Low UV flux (IRAS 05341)

– Significant excess

hydrogen

– -CH3 or -H

M. P. Bernstein, et al., ApJ, 472, L127 (1996).

Reactions out of mass 300

σ = 0.6 ± 0.3 Å2

H-PAH formation

CDH groups

CD2 groups

Subsequent H irradiation

IR spectroscopy

Menella et al., Astrophys. J. Lett. 2012

Experiments reveal many different

pathways to H2 formation, which

together make H2 formation

efficient under the many different

physical and chemical conditions in

the ISM

Interstellar surface reactions

H O

C N

H2

ISM in the Milkyway

Overview picture

ISM and star formation

Eagle-Nebula -Pillars

Eagle-Nebula