Transcript How to observe the H2 component?

Molecules in galaxies
1. How to observe the H2 component?
2. Molecular component of the Milky Way
3. H2 in external galaxies
4. Molecules in absorption
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How to observe the H2
component?
IRAM Summer School Lecture 1
Françoise COMBES
The H2 molecule
• Symmetrical, no dipole
• Quadrupolar transitions ΔJ = +2
• Light molecule => low inertial
moment and high energy levels
• Para (even J) and ortho (odd J)
molecules (behave as two different
species)
512K
170K
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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 H++H2
• n(O)/n(P) ~ exp(-170/T)
• Anormal ratios observed (ISO)
• IR lines J=2-1 at 42 μ, 1-0 at 84 μ ?
• A = 10-10 cm3/s (Black & Dalgarno 1976)
• A=2 10-10 cm3/s (Gerlich 1990) reaction favors o-H2
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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, v=1-0 S(1)
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excitation by shocks, SN, outflows
or UV-pumping in starbursts, X-ray, AGN
require T > 2000K, nH2 > 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)
S(0) filled; S(1) open – CO profile (full line)
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NGC 891, Pure rotational H2 lines S(0) & S(1)
S(0) wider: more extended?
Derived N(H2)/N(HI) = 20 ; Dark Matter?
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H2 v=1-0 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 (Av < 1.5)
• 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) in H2
• (not in CO)
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Ly 4-0
FUSE Spectrum of the LMC star Sk-67-166 (Tumlinson et al 02)
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NH2 = 5.5 1015cm-2
R0
R0/3
Io
Io*20
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
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 ~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
• J=1 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 (J=1-0 at 2.6mm, J=2-1 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 << 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
Mvt=39LCO.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 n~200cm-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) α (O/H)2.2
Confirmed by Taylor & Kobulnicky (98)
But see Walter et al (2003)
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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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High Density Tracers
Nuclei of Galaxies possess denser gas
GMC to survive to tidal forces must be denser
High-J levels of CO
higher critical density to be excited (105cm-3)
as well as
HCN, HCO+, CS, CH3OH, H2CO, OCS, etc..
SiO traces shocks (for instance supershells in starbursts)
Isotopic studies: primary or secondary elements can trace the age
of the star formation events
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M82
Mao et al (2000)
High J-levels of CO
Images are roughly similar in morphology
although somewhat less extended than CO(1-0)
Two hot spots on either side of the nucleus
Part of the molecular torus seen edge-on
Ring due to the bar (or also void due to starburst?)
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High density tracers, at low temperatures
CS, HCN
The ratios CS/CO and HCN/CO are correlated with LFIR
(1/6 in ULIRGs, 1/80 normal, as MW)
Starbursts have a larger fraction of dense gas
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Isotopic molecules
12C/13C
in the MW, from 50-90 at the Sun radius
towards 10-20 in the center
Tracer of the astration, 13C is secondary
In the Galactic Center, also deficiency of deuterium
In Starbursts and ULIRGS (Arp220 type), CO/13CO larger
Not due to a low optical depth, since C18O is normal with
respect to 12CO
But 12C is overproduced in the nucleosynthesis of a recent burst
(Casoli et al 1992)
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12C/13C
ratios determined in M82 and IC342 by Henkel et al 98
from CN, HCN, HCO+ observations
Always
16O/18O
12C/13C
>40 (not as low as in the Galactic Center)
> 100, 14N/15N > 100
HC15N detected in LMC and N4945 (Chin et al 99)
14N/15N = 111
lower than in the Milky Way
==>15N is synthesized by massive stars
Controversial about this formation: destruction in H-burning
formation in SN-II, 14N more secondary, and the ratio increase
with time and astration
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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
Proportional 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 < 1 in the disk of galaxies
Tex (21) < Tex (10) upper level not populated
even if Tkin would have allowed them
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M82, CO3-2, 2-1, 1-0
_|¯ __ - Isotopic ratio of about
10-15 for 12CO/13CO
==> Opticall thick gas
TA* = (Tex -Tbg) (1 - e-τ)
If optically thin R(21/10) --> 4
Survey of CO(3-2) in 30 spiral galaxies (Mauersberger et al 99)
R(32/10)= 0.2-0.7, predicted if Tkin < 50K and n(H2) < 103cm-3
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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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CI/CO = 0.2 (in Kkm/s)
cooling comparable
Normally smaller than C+
(except Arp220 and Mkn231!)
Gerin & Phillips432000
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), and excitation
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
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