physics and chemistry of the diffuse interstellar mediumsemenov/lectures/... · • interstellar...
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Physics and Chemistry of the Diffuse Interstellar Medium
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What is the Diffuse Interstellar Medium?
Picture from Dopita / Sutherland, “Astrophysics of the Diffuse Universe”, Springer
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HII regions
• in Astronomy “HII” stands for ionized hydrogen atoms H+, “HI” stands for neutral hydrogen H, “FeIV” would stand for Fe3+, etc …
• consequently, HII regions are characterized by all hydrogen being ionized,
• Temperatures are on the order of 10 000 K,
• HII regions are low density clouds of gas where star formation recently happened,
• too hot and violent for molecules to exist,
• Precursors are Giant Molecular Clouds -> star formation
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The Cosmic Chemistry Cycle
Credit: NRAO
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Spectral classes of stars
Young, massive stars shine in the UV
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HII regions and the attenuation of Extreme UV light (EUV)
Star
H+
He
UV photons
Photons with 𝐸𝛾 > 13.6 eV directly photoionize H
atoms H + 𝛾 H+ + e- (+ ΔE)
e-
Recombination between electrons and H+ can occur into any atomic level, followed by photon emission:
H+ + e- H(n)
Except for a direct capture into n=1, the emitted photon has an energy of < 13.6 eV, insufficient to ionize neutral H, therefore the UV radiation field is weakened
H(n) H(n’) + 𝛾 (n’ < n)
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O eV
-13.6 eV n=1
n=2
n=3 n=4 n=5 n=6
n ∞
Lα -
12
1.7
nm
Hα
– 6
56
.3 n
m
n=3 n=2 @ 656.3 nm responsible for red glow in H II regions containing ionized H atoms
Lym
an
Bal
mer
Pasc
hen
-3.4 eV
-1.9 eV
HII regions and the attenuation of Extreme UV light (EUV)
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Transition between HII and HI region
From A.G.G.M. Tielens: “Interstellar Medium”, Cambridge press
HII region HI region
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Star
H+
EUV photons
e-
Neutral H
Photons with E < 13.6 eV
HII region
HI region
HII regions around young stars act as filters for EUV photons with E > 13.6 eV
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For completeness: the Intergalactic Medium
Very dilute, hydrogen fully ionized, Very hot: 104-107 K.
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What is the Diffuse Interstellar Medium?
Picture from Dopita / Sutherland, “Astrophysics of the Diffuse Universe”, Springer
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Diffuse medium vs. terrestrial conditions
Earth’s atmosphere Interstellar Medium
Density 1019 cm-3 102 – 106 cm-3
Collision time scale
nanoseconds, 3-body collisions likely
hours – years for binary collisions, no 3-body collisions
statistics
Boltzmann Often local thermodynamic
equilibrium is not reached, radiative lifetimes can be much shorter than
average collision times
many species in ground state, irrespective of kinetic temperature
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Traditional Classification of the diffuse ISM
HIM: hot ionized medium
WIM: warm ionized medium
WNM: warm neutral medium
CNM: cold neutral medium
x: ionization fraction McKee & Ostriker, ApJ 218, 148 (1977)
The ISM is complex in it’s structure, probably clumpy, therefore most sightlines probably consist of a mixture of different phases.
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Modern Classification of Interstellar Cloud Types
Snow & McCall, Annu. Rev. Astron. Astrophys. 2006. 44:367-414
Definitions: Number density of species X: n(X) [cm-3] Column density: N(X) [cm-2] Total nuclei number density: nX [cm-3] Local fraction: fn
Integral of n(X) along sightline
example: NH=n(H)+ 2n(H2), nC ≈ n(C+) + n(C) + n(CO)
example: fnH2= 2n(H2)/nH , f
nCO = n(CO)/nC , etc …
The diffuse ISM can be explored by Visible and UV observations! Easy to study? Still a lot of problems …
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Transition between Cloud Types
Snow & McCall, Annu. Rev. Astron. Astrophys. 2006. 44:367-414
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Diffuse Atomic Clouds
• Part of the diffuse interstellar medium that is fully exposed to radiation field • Nearly all molecules quickly destroyed by photodissociation • Hydrogen in neutral atomic form (H). • Elements with Ionization potential < 13.6 eV almost fully ionized (e.g. C -> C+)
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Diffuse Interstellar Bands: Dependence on N(H2)
• There seems to be no correlation between N(H2) and the strength of the DIB features. • DIBS exist even when N(H2) gets very small. • This suggests that H2 is not involved when the DIB carriers are formed!
Herbig ApJ 407, 142 (1993)
DIBS
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Herbig ApJ 407, 142 (1993)
Diffuse Interstellar Bands: Dependence on N(H)
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Diffuse Atomic Clouds
• Part of the diffuse interstellar medium that is fully exposed to radiation field • Nearly all molecules quickly destroyed by photodissociation • Hydrogen in neutral atomic form (H). • Elements with Ionization potential < 13.6 eV almost fully ionized (e.g. C -> C+) • Hardly any molecules, but strong Diffuse Interstellar Bands (Mystery!)
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Diffuse Molecular Clouds
• Interstellar radiation field is sufficiently attenuated in the UV regime for H2 to survive,
• typically surrounded by diffuse atomic gas, • Most diffuse molecular sightlines will also cross atomic sightlines, • densities 100-500 cm-3, • temperature 30-100 K, • Almost all carbon ionized (C+) • High ionization fraction, n(C+) ≈ n(e) • Electron recombination dominant destruction route for many ions
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Translucent Clouds
• Characterized by the transition from ionized C+ to neutral C and CO • Introduced by Van Dishoeck and Black (1989), • Translucent regime not well understood, • Lack of observational data, • Chemical models do not agree in the C, Co transition regime,
• Interstellar radiation field is sufficiently attenuated at energies <13.6 eV
that carbon becomes neutral (C) or molecular (CO) • Typically surrounded by diffuse molecular gas,
Translucent cloud
Diffuse Molecular Cloud (stops UV with Eγ >11eV)
Diffuse atomic Region (stops EUV
with Eγ > 13.6 eV )
H
H2
C+
H2
C
CO
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Dense Molecular Clouds
• With increasing extinction, carbon becomes completely molecular (CO), • No longer observable in the visible or UV regime, • Much lower ionization fraction than diffuse medium, • Electron recombination is much less important due to lack of electrons, • Main source of ionization: Cosmic rays! • Most of the detected interstellar molecules are observed in dense clouds, • Places of star formation.
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Key Species in Diffuse Molecular Clouds: H2
• No dipole moment, homonuclear no allowed vibrational or rotational transitions
• Space-based far-UV spectrographs needed for observations
Lyman bands: Eγ >11.2 eV λ < 110.8 nm
Werner bands Eγ >12.3 eV λ < 100.8 nm
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First Interstellar H2 Absorption spectrum
• Rocket experiment : Carruthers, ApJ 1970, 161 L81 • H2 and H column densities comparable
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Far Ultraviolet Spectroscopic Explorer
FUSE 1999-2007
H2 UV observations with FUSE
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Observed H2 rotational population
Rotational excitation
𝑁𝐽 ~ 𝑔𝐽 𝒆 𝑬𝑱
𝒌𝑻 Boltzmann law:
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Key Species in diffuse clouds: H2 Formation
H + H H2 + γ
Direct radiative association:
The H2+ pathway:
H + H+ H2+ + γ
H2+ + H H2 + H+
forbidden
Radiative transitions between the lowest states of H2 forbidden, fraction of atoms in 2p states negligible
Radiative association extremely inefficient
𝛼 < 10−16 cm3 s-1
𝛼 ≈ 10−9 cm3 s-1
The 1. step is rate-limiting, renders H2+ path negligible
The H- pathway:
H + e- H- + γ H- + H H2 + e-
Without shielding of visible light, H- does not live long enough
H- binding energy 0.75 eV
Only formation on dust grains remains
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H2 formation on grains
E. Van Dishoeck / Lecture notes
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H2 observations vs. dust grains
Van Hoof et al., A&A 2010
H2 2.12 μm Ground based Carlo Alto Herschel
(NGC 6720 plantery nebula)
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Ro-vibrational excitation in HD formation detected by Resonantly Enhanced Multiphoton Ionization (REMPI)
F. Islam, Thesis 2009 University College London
multiphoton excitation
ionization
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HD Formation and Ro-vibrational excitation
HD formed in v=4
Trot= 368 +- 22 K
Energy released in the Dust grain formation may contribute to the non-thermal H2 distribution observed in diffuse clouds
F. Islam, Thesis 2009 University College London
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H2 optical pumping
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Observed H2 rotational population
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Rotational excitation of H2
kinetic temperature and optical pumping H2 lines up to J = 7 detected with Copernicus + FUSE Population distribution non-thermal, 2 temperature regimes! Low J: excitation dominated by collisions • sensitive to T and nH • abundance H+ large enough that ortho/para exchange rapid
and J=1/J=0 gives Tkin
High J: energy levels lie very high (> 1000 K) • not populated by collisions at T = 40 - 80 K • populated by optical pumping • process through B <- X and C <- X systems • proportional to radiation field • possibly contribution from dust grain formation process
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H2 nuclear spin and kinetic cloud temperature
Rotational excitation
ΔE (J=0, J=1) = 170.9 K
Ortho – para exchange reactions
H + o-H2 H + p-H2
Activation barrier too high
0.5 eV / 6000 K
H+ + o-H2 H+ + p-H2
• No activation barrier • Requires some ionized atomic H fraction
1)
2)
Due to high efficiency of 2), H2 ortho / para fraction should reflect kinetic cloud temperature
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H2 destruction: possible processes
For direct photodissociation photons with Eγ > 15 eV are needed, which are already Depleted by the atomic outer layers
There is no bound state in the right energy regime for H2
Major H2 destruction mechanism in the diffuse ISM
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H2 destruction by radiative dissociation through Lyman and Werner Bands
90% of the flux returns back to the electronic ground state (contributing to the observed optical pumping of rotational states)
10% of the flux leads to spontaneous dissociation
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Depth dependence of various species
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HD observations in the diffuse medium
FUSE UV observations
• Even most HD lines are fully saturated • Obscured by H2 and other molecular lines • Difficult analysis
Can we use N(HD) to constrain cosmic D/H ?
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N (HD) does not seem to represent cosmic H/D fraction (≈ 10-5 ). Why?
HD fraction in diffuse clouds
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Processes that influence the HD fraction
• H2 is protected by self shielding of narrow resonances • HD self-shielding sets in much later
H2 is favored, low HD/H2 ratio
• D atoms have a lower mobility on grains, formation mechanism may be much slower
• HD has a gas-phase formation channel H2 + D+ HD + H+ (exothermic)
(depends on ionization rate)
Too complicated to infer reliable Information on D/H.
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Summary H2 / HD
• Observed in the diffuse ISM through UV satellite observations, • Formed exclusively on dust grains, • Ortho/para ratio reflects the cloud kinetic temperature (50-100 K), • Destroyed through radiative dissociation via Werner and Lyman bands, • Non-thermal rotational populations in the outer layers (optical pumping), • Increasing self-shielding with cloud depth.
H2
HD
• Can also be observed in the diffuse ISM through UV satellite observations, • Has a gas-phase formation rate, • reduced self-shielding pushes N(HD)/(2 N(H2)) below cosmic D/H.
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PAUSE
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Other interesting molecules: CN and the CMB
• One of the first interstellar molecules that were discovered (McKellar PASP 1940) (actually Interstellar molecule no 2, after CH in 1937)
• Large dipole moment -> in low-collision environment, state population reflects radiation field
• Rotational temperature of (2.3 ± 0.5) K derived by McKellar in 1940 pre-discovery of the CMB
• Origin of this temperature was not understood until 1965
Anecdote: Gerhard Herzberg, in his textbook on Diatomic Molecules 1950 noted that the 2.3 K rotational temperature of the cyanogen molecule (CN) in interstellar space is interesting, but stated that it had "only a very restricted meaning.“ Missed Nobel Prize in Physics?
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The mysterious abundance of CH+
• Discovered in 1941 (Herzberg) • First interstellar molecular ion • Thought to be created by
C+ + H2 → CH+ + H (endothermic by 4600K)
Where does the energy come from to overcome the reaction barrier?
Candidates: Shocks? Magnetohydrodynamic waves? Alv’en waves? Turbulence?
Why do these mechanisms not enrich other endothermic species?
Presently no answer.
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Plasil et al., ApJ 737, 60 (2011)
CH+ + H C+ + H2
The rate coefficient for the reaction
is a factor of 50 smaller than the classical Langevin value at low energy
The mysterious abundance of CH+
In equilibrium the formation and destruction processes will determine the abundance of CH+
New experiment (2011): CH+ + H C+ + H2
22 pole ion trap
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H3+: the Engine of Interstellar Chemistry
temperatures: 10-100 K density: 102 - 104 cm-3
Molecular clouds: It’s cold and (relatively empty)
Key for active low-temp chemistry: Ion-neutral Reactions!
• No endothermic reactions • No 3-body collisions • No reactions with barriers
H2+
H3+
CH+
CH2+
CH3+
CH5+
CH4
C2H3+
C2H2
C3H+
C3H3+
C4H2+
C4H3+
C6H5+
C6H7+ C6H6
H2
H2
H2
H2
H2
C
e
C+
e
C+
C
H
C2H2
H2 e
OH+ H2O+
H3O+ H2O
OH e
O
H2
H2
HCO+ CO
HCN CH3NH2
CH3CN
C2H5CN
CH CH2CO
CH3OH
CH3OCH3
CH3+
C2H5+ e
C2H4
e
C3H2 e
C3H
e
C2H
McCall, PhD thesis, Chicago 2001
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First quantitative cloud model (Herbst & Klemperer 1973)
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Extraterrestrial H3+
1.02
1.00
0.98
0.96
0.94
36700366803666036640 3717037150
AFGL 2136
AFGL 2591
Geballe & Oka, Nature 384, 334 (1996)
Dense clouds
McCall et al., Science 279, 1910 (1998)
Diffuse clouds
Connerney et al., Science 262 , 1035 (1993)
Jupiter
Plus: • detection in the atmosphere of Uranus • detection in Saturn’s atmosphere • high abundance towards the galactic center • extragalactic detection in IRAS 08572+3915NW
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Extraterrestrial H3+
Jupiter: The Planet, Satellites and Magnetosphere, Ed. F. Bagenal, T. Dowling and W. McKinnon Cambridge University Press (2004)
H3+ auroral emission
from the Jovian atmosphere
Connery and Satoh, Phil. Trans. R. Soc. Lond. A, 358, 2471 (2000)
Plus: • detection in the atmosphere of Uranus, • detection in Saturn’s atmosphere, • high abundance towards the galactic center.
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The enigma related to interstellar H3+
Dense clouds: formation H2 + c.r. H2+
H2+ + H2 H3
+ + H destruction H3
+ + CO HCO+ + H2
observed column density: 6 x 1014 cm-2
Diffuse clouds: formation H2 + c.r. H2+
H2+ + H2 H3
+ + H destruction H3
+ + e- H2 + H / H + H + H
observed column density: 4 x 1014 cm-2
2-3 orders of magnitude too much H3+ in diffuse clouds !
Dissociative Recombination DR
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H3+
H2 + H H + H + H
e-
The Problem with H3+ Electron Recombination
H3
+ D
R r
ate
coef
fici
ent
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Steady State: [H3+]
=
ke [e-]
[H2]
The Enigma of H3+ in Diffuse Clouds
Cosmic ray ionization rate (10-17 s-1)
Fast DR rate coefficient (10-7 cm3 s-1)
With the canonical value for the Cosmic Ray Ionization Rate Of ≈ 10-17 s-1 and a fast DR process,
H3+ should not be observable in the diffuse ISM!
Inverse ionization fraction (1/10-4)
N(H3+) = 10-6 cm-3
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Why H3+ observations are so useful
• H3+ density n(H3
+) is constant with two phases in diffuse and dense sightlines • By measuring the column density, the length of the cloud L along the line
of sight can be calculated:
L = N(H3+)/n(H3
+)
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Storage Ring Recombination Measurements
• radiative relaxation (rotations, vibrations) • direct measurement
• 100% detection efficiency
• high resolution
Advantages
Problem
H3+
vibrational
excitation?
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H3+ Vibrational Cooling
first breathing mode (1,xx)
(1,00) theory
(2,00) theory
experimental data (T > 2s)
theory ground state (0,00)
experimental data (T < 1ms)
Kreckel et al., PRA 66, 052509 (2002)
Rotational Excitation 3000 K !
Kreckel et al., New J. Phys. 6, 151 (2004)
Metastable rotational states
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H2 Gas inlet 2 atm
Solenoid valve
-900 V ring
electrode
Pinhole flange/ground
electrode
H3+
Rotationally “cold” Ion Sources
Supersonic expansion
CRYRING Stockholm
H3+
Cryogenic 22-pole ion trap
TSR Heidelberg
Gerlich, Physica Scripta, T59, 256, (1995)
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H3+ DR Spectrum High Resolution
● TSR electron temp. kT┴=0.5 meV ~ 6K kT||=25 μeV trap temp. 13 K Kreckel, PRL 95, 263201 (2005)
● CRYRING kT┴= 2 meV ion temp. 30-50 K McCall, Nature 422, 500 (2003)
Theory: Kokoouline and Greene
Nature (News and Views) 440, 157 (2006)
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Steady State: [H3+]
=
ke [e-]
[H2]
Cosmic ray ionization rate
Fast DR rate coefficient
High observed H3+ column densities
High
H3+ Laboratory Astrophysics Research leads to a revised
Cosmic ray ionization rate in the Diffuse ISM
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Nuclear Spin equilibrium in the diffuse ISM
Crabtree et al., ApJ 729, 15(2011)
T(H2) ~ 60 K
T(H3+) ~ 30 K
10 Thermal Distribution
• Astronomical Observations
Solution:
H3+ + H2 → (H5
+)* → H3+ + H2
or
H3+ + e-
H3+ colder than H2 !
Too much para-H3+
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T(H2) ~ 60 K
Measure with thermal p-H2 samples in 10k steps
n-H2
100% p-H2
90% p-H2
p p o o Raman-Spec.
Understanding the H3+ - H2 collision system
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Detected Molecules
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Processes in Diffuse Clouds
Ewine van Dishoeck, lecture notes
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Carbon / Oxygen Network
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Carbon / Nitrogen Network
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Depletion on Dust Grains: also in the diffuse medium
More depletion in cool clouds, more freeze-out on dust grains
Savage & Sembach, Annu. Rev. Astron. Astrophys. 1996. 34:279-329
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Elemental Depletion
𝐷𝑥 = log𝑁 𝑥
𝑁 𝐻− log
𝑁(𝑥)
𝑁(𝐻)⊙
Timescale for depletion:
𝜏𝑑 =1
𝜀 𝑓 𝑋 𝑣 𝑋 𝑛
ε Grain surface per H atom cm-3
𝑓 𝑋 Sticking probability for species 𝑋
𝑣 𝑋 velocity of species 𝑋
𝑛 cloud density
Typical cloud lifetime: 3 x 1014 s
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Non-detections
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Some notable reaction chains
H3+ formation
and destruction H2 H2
+ H3+
cosmic ray H2 H2 + H
H + H + H
e-
O H3
+
OH+ OH2+
H2 OH3
+ H2 e-
OH + H2
H2O + H
Water formation
C+ OH
CO+ H2
HCO+ e-
H + CO Carbon chemistry
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Herschel surprisingly finds a lot of H2O+
Regions with low H2 density?
O H3
+
OH+ OH2+
H2 OH3
+ H2 e- OH + H2
H2O + H
Water formation
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Or: maybe it’s more complicated
J. Stutzki, “On the Complex Structure of Interstellar Clouds”, 2009
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Summary • The diffuse interstellar medium can be classified into 3 phases:
1. diffuse atomic clouds (hydrogen atomic H) 2. diffuse molecular clouds (hydrogen molecular H2, carbon ionized C+) 3. translucent clouds (hydrogen molecular H2, carbon neutral C)
• Active Ion-neutral collisions build up larger molecules, • High electron fraction -> Electron recombination important, • H2 becomes self-shielding, • H2 ortho/para excitation temp. reflects the temperature of the cloud (10-100 K), • HD can be observed easily, exact abundance difficult to predict, Mysteries Remain: • Formation of CH+
• H3+ column densities imply higher
cosmic ray ionization rate, enigma solved? • Why so much H2O+? • Why is H3
+ colder than H2?
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Literature
Snow & McCall, “Diffuse Atomic and Molecular Clouds” Annu. Rev. Astron. Astrophys. 2006. 44:367-414
Tielens, A.G.G.M. “The physics and chemistry of the interstellar medium” Cambridge University Press