How to observe the H2 component?

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Molecules in galaxies at all redshifts 1.
How to observe the H2 component? 2. Molecular
component of the Milky Way 3. Fractal
Structure 4. Formation of the fractal, shear,
turbulence 5. H2 in external galaxies 6. H2 in
ULIRGs, Dense tracers 7. Molecules in
absorption 8. CO at high redshift 9. Primordial
H2, history
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How to observe the H2 component?

• SAAS-FEE Lecture 1
• Françoise COMBES

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The H2 molecule

• Symmetrical, no dipole
• Quadrupolar transitions ?J 2
• Light molecule gt low inertial moment and high
  energy levels
• Para (even J) and ortho (odd J) molecules (behave
  as two different species)

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H2 is the most stable form of hydrogen at low
T dominant in planetary atmospheres? Formation
on dust grains at 10K However formation still
possible in primordial gas (H H- Palla et al
1983) Destruction through UV photons (Ly
band) Shielded by HI, since the photodissociation
continuum starts at 14.7eV, and photo-ionization
at 15.6 eV (HI ionization at 13.6
eV) Self-shielding from low column
densities 1020 cm-2 in standard UV field H2 will
be present, while other molecules such as CO
would be already photo-dissociated
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Potential curves involved in the Lyman and Werner
bands (Roueff 00)
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Ortho-Para transitions?

• Formation in the para state not obvious
• Large energy of formation 2.25 eV/atom
• ortho-para conversion in collisions HH2
• n(O)/n(P) 9.35 exp(-170/T)
• Anormal ratios observed (ISO)
• IR lines J2-1 at 42 µ, 1-0 at 84 µ ?
• A 10-10 cm3/s (Black Dalgarno 1976)

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Infrared Lines of H2

• Ground state, with ISO (28, 17, 12, 9µ)
• S(0), S(1), S(2), S(3)
• From the ground, 2.2 µ, v1-0 S(1)
• excitation by shocks, SN, outflows
• or UV-pumping in starbursts, X-ray, AGN
• require T gt 2000K, nH2 gt 104cm-3
• exceptional merger N6240 0.01 of L in the 2.2 µ
  line (all vib lines 0.1?)

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H2 distribution in NGC891 (Valentijn, van der
Werf 1999)
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NGC 891, Pure rotational H2 lines S(0) S(1)
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H2 v1-0 S(1) 2.15µ in NGC 6240 van der Werf et
al (2000) HST
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UV Lines of H2

• Absorption lines with FUSE
• Very sensitive technique, down to column
  densities of NH2 1014 cm-2
• Ubiquitous H2 in our Galaxy (Shull et al 2000,
  Rachford et al 2001) translucent or diffuse
  clouds
• Absorption in LMC/SMC reduced H2 abundances, high
  UV field (Tumlinson et al 2002)
• High Velocity Clouds detected (Richter et al 2001)

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Ly 4-0
FUSE Spectrum of the LMC star Sk-67-166
(Tumlinson et al 02) NH2 5.5
1015cm-2
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R0/3
R0
Io
Io20
Column densities and molecular fraction compared
to models
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Detection of H2 in absorption by FUSE in HVCs
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Sembach et al 2001
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The CO Tracer

• In galaxies, H2 is traced by the CO rotational
  lines
• CO/H2 10-5
• CO are excited by collision with H2
• The dipole moment of CO is relatively weak
• ? 0.1 Debye
• Spontaneous de-excitation rate Aul ? ?2
• Aul is low, molecules remain excited in
  low-density region about 300 cm-3

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• Competition between collisional excitation and
  radiative transitions, to be excited above the
  2.7K background
• J1 level of CO is at 5.2K
• The competition is quantified by the ratio
  Cul/Aul
• varies as n(H2)T1/2 /( ?3 ?2)
• Critical density ncrit for which Cul/Aul 1
• Molecule CO NH3 CS HCN
• ? (Debye) 0.1 1.5 2.0 3.0
• ncrit (cm-3) 4E4 1.1E5 1.1E6 1.6E7

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Various tracers can be used, CO for the wide
scale more diffuse and extended medium, the dense
cores by HCN, CS, etc.. The CO lines (J1-0 at
2.6mm, J2-1 at 1.3mm) are most often optically
thick At least locally every molecular cloud is
optically thick Although the "macroscopic" depth
is not realised in general, due to velocity
gradients Relation between CO integrated
emission and H2 column density? Is it
proportional? How to calibrate?
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NGC 6946 CO(2-1) map 13" beam IRAM 30m Spectra,
Weliachew et al 1988
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? Isotopic molecule 13CO, UV lines ? Statistics
of "standard" clouds ? The Virial relation 1-
Use the isotope 13CO much less abundant at the
solar radius Ratio 90 therefore 13CO lines
more optically thin A standard cloud in the MW
has ?CO 10 and ?13 0.1 The average ratio
between integrated CO and 13CO intensities is of
the order of 10
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Successive calibrations knowing 13CO/H2 ratio in
the solar neighbourhood (direct observations of
these lines in UV absorption in front of stars,
with diffuse gas on the line of sight) 2-
Statistically "standard" clouds For extragalactic
studies, numerous clouds in the beam Typical
mass of a cloud 103 Mo something like 104 or 105
clouds in the beam No overlap, since they are
separated in velocity Filling factor fs fv ltlt 1
(hypothesis) Usually TA 0.1K for nearby
galaxies, 10K for a cloud
constant factor between ICO and NH2
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3- More justified method the virial Each cloud
contributes to the same TA in average reflecting
the excitation temperature of the gas the width
of the spectrum gives the cloud mass through the
virial hypothesis V2 r GM The conversion ratio
can then be computed as a function of average
brightness TR and average density of clouds n
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Milky Way Virial mass versus LCO Mvt39LCO.81
Slope is not 1
Solomon et al 1987
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Area A of the beam A ?/4 (?D)2 N clouds, of
diameter d, projected area a ?/4 d2 velocity
dispersion ?V ICO A-1 N (?/4 d2)TR ?V Mean
surface density NH2 A-1 N (?/6 d3) n NH2 /
ICO 2/3 nd/( TR ?V) from the Virial ?V n1/2
d and the conversion ratio as n1/2 /TR
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This factor is about 2.8 E20 cm-2 /(km/s) for TR
10K and n200cm-3 This simple model expects a
low dependence on metallicity, since the clouds
have high optical thickness and are considered to
have top-hat profiles (no changes of sizes with
metallicity) However, for deficient galaxies
such as LMC, SMC, where clouds can be resolved,
and the virial individually applied, the
conversion factor appears very dependent on
metallicity
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The size of clouds, where ? 1, is varying
strongly Models with ? r-2, NH2 r-1 Diameter
of clouds d Z (or O/H) Then filling factor in
Z2 The dependence of the conversion ratio on
metallicity could be more rapid than linear (the
more so that C/O O/H in galaxies, and CO/H2
(O/H)2) In external galaxies, the MH2/MHI
appears to vary indeed as (O/H)2 (Arnault et al
88, Taylor et al 1998)
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Arnault, Kunth, Casoli Combes 1988
LCO/M(HI) a (O/H)2.2
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On the contrary, in the very center of starbursts
galaxies, an overabundance of CO could
overestimate the molecular content Not clear and
definite variations, since TR is larger, but nH2
too, and NH2 / ICO varies as n1/2 /TR Possible
chemical peculiarities in starbursts 12C primary
element, while 13C secondary Isotopic ratios vary
Can be seen through C18O
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Another tracer cold dust
At 1mm, the emission is Rayleigh-Jeans B(?, T)
2 k T / ?2 flux quasi-linear in T (between 20 and
40K) In general optically thin emission Proportio
nal to metallicity Z Z decreases exponentially
with radius
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When the molecular component dominates in
galaxies, the CO emission profile follows the
dust profile (example NGC 891) When the HI
dominates, on the contrary, the dust does not
fall as rapidly as CO with radius, but follows
more the HI (example NGC 4565)
CO might be a poor tracer of H2
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Radial profiles N891 (Guélin et al 93) N4565
(Neininger et al 96)
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The excitation effects combine to
metallicity Explains why it drops more rapidly
than dust with radius CO(2-1) line tells us
about excitation Boarder of galaxies, CO
subthermally excited When optically thick
CO21/CO10 ratio 1 If optically thin, and same
Tex, could reach 4 But in general lt 1 in the disk
of galaxies Tex (21) lt Tex (10) upper level not
populated even if Tkin would have allowed them
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Braine Combes 1992, IRAM Survey
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Gradient of excitation in the LMC vs MW Sorai et
al (2001) Average value of 0.6 for MW
from Sakamoto et al 1995
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CO(2-1)/CO(1-0) vs IRAS, and vs CII in LMC (grey
band MW)
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Conclusion
The H2 molecule is invisible, in cold molecular
clouds (the bulk of the mass!) CO is not a good
tracer, both because metallicity effect (non
-linear, since depending on UV flux,
self-shielding, etc. Very important to have
other tracers dense core tracers, HCN, HCO,
isotopes.. H2 pure rotational lines, also a
tracer of the "warm" H2, always present when cold
H2 is there