History of the molecular component

SAAS-FEE Lecture 9

Françoise COMBES

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Formation of H2

Formation of the H2 molecules, very important in the early universeMain coolant, to allow the collapse of structures (then HD..)

at z=0, the dust acts as a catalyst

As soon as dust is present, and ions have decreased, H2 is again formedon dust

Formation on grains: two processes:Langmuir-Hinshelwood, H atoms land on a grain and diffuse on the surface, by either tunneling, or hopping

Eley-Rideal, H atoms are fixed in position, chemisorbed,and reaction occurs only when another H lands atop

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Binding sites for adsorbers: N=106 on a 0.1μ grain

Binding energy varies: H on olivine ED = 372K (Katz et al 99)

H on amorphous carbon ED = 658K For H2 respectively 314 and 542K (Katz et al 99)

Since the desorption energy is larger for heavier molecules, onlyH2 can be formed like that, at 10K

The grain, as a catalyst, absorbs the energy of formation, that canhelp to desorb another molecule

At low temperature; H-atoms dont diffuse on the grain; and the second(Eley-Rideal) process is the only one

==> Still poorly understood (Herbst 2000, « H2 in Space  » CUP)

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Formation on Coronene (PAH) Sidis et al 2000

Computations of Hwith C24H12

Formation of H- on MgOor forsterite Mg2 SiO4

by charge transferAb initio

Experimental formation of H2 on amorphous Cor olivine (Pirronello et al 1999)

Recombination efficiency versus Ttriangles: during irradiation (48s)squares: during desorptioncircles: total

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Formation in gas-phase reactionsLepp & Shull 1984

Formation of the interesting molecules H2, HD, LiH and HeH+during the post-recombination epoch (z=300-30)

In standard nucleosynthesis, there exist only H, D, 4He, 3He, 7Li

Cooling by Ly is ineffective below T= 8000K, after z=500

H2 is formed in small quantities from H-, and provides sufficient cooling to trigger the cloud collapse

Cooling is necessary since a Jeans mass will increasein adiabatic collapse==> no fragmentation

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First Process (Bienik & Dalgarno 79)H + e- --> H- +

H-+H --> H2 + e-

Second Process (Karpas et al 79)

H+H+ --> H2+ +

H2+ + H --> H2 + H+

Third Process (Palla et al 83)Three-body process, at high density 109 cm-3

3H --> H2 + H

Direct radiative association forbidden, since no dipole of H2

Not the case for HD and LiH (H+H --> H2 + )

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Results of chemistry

Standard FriedmanΩ0= Ωb = 0.1

Channels indicated

Once Compton heatingbecomes less than expansion cooling,radiation and matter cool adiabatically

Tr ~(1+z) Tm ~ (1+z)2

Lepp & Shull 1984

H2 reaches ~ 10-5

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During cloud collapseCooling rate per H2 molthermal equilibriumLiH dominatesfor HD/H2=10-4

LiH/H2=10-6

n= 3 1010 cm-3

Self-similar free-fall collapseat z=50, λJ = 560pc, MJ=105Mo

More molecules formed

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The cooling is enough for cloud to collapse, without beingstopped by adiabatic regime

The cooling rotational lines become optically thick at

N = 1024 cm-2 Δv for H2

N = 1020 cm-2 Δv for HDN = 1015 cm-2 Δv for LiH

Δv ~100km/s estimated by Lepp & Shull 1984H2 becomes optically thick at the same point as 3-body reaction sets in F = 1012

LiH remains thin and effective coolant until higher densitiesLiH/H = 10-10

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Chemistry of early UniverseMore detailed reactions (Galli & Palla 1998): 87 gas phase reactions

H, D, He, Li chemistry

Different: radiative associationfor LiH was overestimatedby 100 (Stancil et al 96)

Non equilibrium chemistryin radiative shocks duringcollapse, etc..

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The minimal network ofchemical reactions, which can reproduce the more complexnetwork

(Galli & Palla 98)

Improvement: better Hrecombination2-3 less electrons

at z=1, e/H = 3 10-4

H2/H ~10-6

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Cooling functions of H2 for H-H2 collnh = 0.1 and 106 cm-3

LS: Lepp & Shull 83HMK Hollenbach & McKee 89MSM Martin et al 96FBDL Forrey et al 97

Cooling function per molecule of H2, HD, LiH and H2+ in the low n (<100cm-3) limit

Galli & Palla 98

Λ=Λo/(1+nc/n)

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Comparisons betweenmodels

Palla et al 95 Lepp & Shull 84

Puy et al 93 Black 91

Giroux & Shapiro 96

Influence ofcosmology

H2 independentLi more variableLiH and HD can vary by 2-3 orders

Now Ωo ~1

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First structures in the dark age

Abel & Haiman 00

In the current best CDM model, the first structures to formare the smallest (bottom-up)

|δk|2 =P(k) ~ kn, with n=1on large-scalesn= -3 on small scalestilt when ρr ~ ρm

at the horizon scale

δM/M ~M-1/2 -n/6

when n > -3, hierarchicalclustering

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Fluctuations can be adiabatic, but they are damped below3 1013 (Ωh2)-5/4 Mo (Silk 1968)or isothermal

The largest masses to get non-linear after recombination(z=1500) are M ~ 106-8 Mo (according to n, and compatible to 1014Mo becoming non-linear today)

Mass spectrum assumingP(k) ~k at large-scaleand P(k) tilted n= -3at small scales (Peebles 82)

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Growth of adiabatic fluctuationsat a scale of 1014Mo (8 Mpc)They grow until their mass equalthat contained in the horizon ctCst after, and calibrated now

Fluctuations of matter (…) ‘standard ’follow radiation, if ionised

After R (recomb) they grow too slowly to reach the observed level

When DM is non-baryonic, only gravity ==> they grow from Ewhen the matter gravity dominates radiation

When HDM, particles decouple when they are relativisticstream motions ==> stabilise

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Two possible forms of spectrumfor fluctuations at Recombination

I: 106-8Mo first, then bottom-up

A: first scales are 1014Mofor HDM (neutrinos): cut-off

I A

Evolution of spherical inhomogeneity

At zsep, decoupling from the rest ofuniverse

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The Jeans mass is M ~105 Mo (Ωb/0.06) -1/2 (h/0.5) -1

This corresponds to Giant Molecular Clouds at z=0

Nature of the collapse

if efficient cooling, quasi isothermaltcool ~ tff (Hoyle 1953)Fragmentation could be pursued until very low masses, in veryshort times (~Myr)

Opacity limited fragmentation

==> M ~ 4 10-3 T 1/4 μ-9/4 f -1/2 Mo

0.1 < f < 1 radiation efficiency (with respect to black-body)

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Mcl ~ 10-3 Mo with densities n ~ 1010cm-3 (PC94)

Is cooling efficient?Compton cooling, HI cooling above 1000K, then H2 cooling

(Yoneyama 1972, Hutchins 76, Carlberg 81, Palla et al. 83, Lepp & Shull 84, Haiman et al 96, Abel et al 97, Tegmark et al 97)

The calculated masses of the fragments rangebetween 0.1 and 100 Mo

More realistic: collapse non spherical, but sheets andfilaments (Larson 1985)

for instance, in 1D pressure forces cannot halt the collapse(Uehara et al 96, Inutsuka & Miyama 97)

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Most studies are interested in star formation

Globular cluster formation (Peebles & Dicke 1968)(very efficient star formation)==> wide mass spectrum of stars (from 0.1 Mo)

With rotation, low-mass as well as VMOs (Kashlinsky & Rees 83)

It is however possible that most of the fragments remain pressure supported, as soon as optically thick

and form a fractal of molecular clouds

All the hierarchical structures remain at the TCMB temperature(PC94)

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H2 formation and cooling

Simple scheme of chemistry reactions (Tegmark et al 97)

H+ + e- --> H + hν (recombination)H + e- --> H- + hνH- + H --> H2 +e-

H+ +H --> H2+ +hν

H2+ +H --> H2 + H+

H+H+H --> H2+HH+H+H2 --> H2+H2

The cooling is essentially due to H2 below 1000KThen HD intervene below 100KLiH has a non-significant abundance (Stancil et al 96)

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Density computed from the simple top-hat model before virialisation

Then assuming efficient cooling, fragmentation occurs, as a fractalstructure D = 1.7N (fragments) =8compatible with the fractal of the Milky Way ISM today(Larson 81, Scalo 85)

==> scales in the ratio N 1/D = 3.4

densities in the ratio rd = N(3-D)/D = 4.9

n(t) = ni rd(t-ti)/tffi

Fragmentation stops when the clumps become opaque to the H2 rotational IR lines, i.e. n ~ 1010 cm-3

Tvir ~1420K (M/105Mo)2/3 (1+zvir)/100When T < Tvir, fragmentation can proceed

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Evolution of the densityinside the bound structures

compared to the backgrounddensity (dash)

The solid lines correspond to the recursive fragmentationwith time-scale as a function of the local density

with different initial fluctuations intensities(Combes & Pfenniger 98)

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n3 density in 103cm-3

fH2 molecular fractionTvir virial temperaturexe ion fraction

T6 gas temperature in 106K

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n3 density in 103cm-3

fH2 molecular fractionTvir virial temperaturexe ion fraction

T6 gas temperature in 106K

Tvir ~(1+zvir) n larger

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n3 density in 103cm-3

fH2 molecular fractionTvir virial temperaturexe ion fraction

T6 gas temperature in 106Kn lower, since dense structuresform first

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What are the first structures

This old question has received many different answers, accordingto assumptions on the IMF and efficiency of star formation

could be gas clumps, brown dwarfs, 0.5 Mo MACHOs, standardIMF stellar systems, or massive black holes

Numerically, often spherical symmetry (Bodenheimer 96)with hydrodynamics, chemistry, (but of course no fragmentation)

Big problem of the high dynamical range (8-15 orders)non-equilibrium chemistry, non -linear dynamics,Problem when the structures become optically thick

Abel et al (95-98): in CDM, blobs form at filaments crossing

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Abel et al 2000

Top: log overdensitycut in the highest peak320pc, 32pc

+ velocities

Bottom: log temperature

Less than 1% of the gaswill form stars

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Baryonic density around the fragmentmass enclosed M( r )compared to Bonnort-Ebert

f(H2) and electron fraction

H2 cooling time, tcross, tff

Temperature (vertical is min)

Virial radius of 5.6 106 Mo halois 100pccell size =0.024pcsimulation box 6.4 kpcAMR

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The first stars or quasars are likely to appear in clumps of Tvir ~104Kat z ~20, and H2 cooling plays a major roleThe structures are probably at filaments crossing

The direct H2 radiation will not be detectable, but as soon as there isstar formation with > 1% efficiency, or quasar radiating at Eddington limit, it will be detectable with NGST at z=15 (Haiman & Loeb 98)

In case of mini-quasars, there will be reionisation by z ~10

The average SF efficiency can be probed through the enrichment ofLyα forestIt is not larger than 2% in collapsed regions, not far from the 1%deduced from numerical simulations

But based on the normal IMF (very uncertain!)

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Efficiency of star formationAlways very low efficiency, since there is not yet a giant potential

At high redshift (z=10-100), the physical conditions of the gas is similar to the present one in the outer parts of galaxies

¶ No deep potential wells, since galaxy-size structures have notyet virialized¶ No critical surface density for star formation ¶ Collapsing mass always of the order of the Jeans mass ==> not far from pressure supportsince the Jeans mass decreases gradually as fragmentation proceeds

When objects of galaxy-size collapse, M >> MJ==> violent instabilities, starburstunless rotation stabilises

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Only sporadic star-formation in small-mass GMC at high z

Can be sufficient to reionize and reheat the diffuse inter-galactic mediumLyα absorbers (Tegmark et al 1984)HeII gas (Jacobsen et al 1994, Davidsen et al 96)

Clumps resistance to reionizationN(H2) ~1025 cm-2 self-shielding from radiation+ fractal structure: more gregarious than homogeneous distribution

Some erosion at interfacesExtragalactic background: estimation of the ionization rate ofξ ~2 10-14 /s (Madau 92)

==> sharp transition (Stromgren sphere theory) in galactic disk atN(HI) ~1018 cm-2 (Corbelli & Salpeter 93)

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Galaxy formation

Could range between z=50 and z=0(at z=200, normal 10kpc galaxies would overlap)

the protogalaxy collapse is violent ( M >> MJ)starburst in their centers

In the outer parts, there can remain cold self-gravitating gasthen hierarchical mergingCluster formation

Another peak of star-formation, when larger structures begin to collapseat z=2, outer gaseous haloes are stripped, heatedFormation of the hot IGM, at Tvir of the clusterMultiphase medium survives (Ferland et al 94)Cooling gas at the center (David et al 95)

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Cooling flows in clusters

The hot X-ray gas pervading the cluster cools and condensestoward the center (generally cD galaxy)Gas phase at intermediate temperature 5 105 K has been seen (EUVE radiation, Lieu et al 96, 99)

Final state of the gas still a mystery, not sufficient stars formed

Recent detection of CO molecules in large amount(Edge et al 01, Salomé & Combes 02)

Cooling filaments extend over kpcMulti-phase gas

Generally associated with starbursts, and also AGN(feedback? Stop the cooling flow?)

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Latest Chandra results on Abell 1795(ACIS for 11h) Fabian et al. 2001size is 75" on a side

Hα + NII (6583 A) emittinggas surrounding the Abell 1795cluster central Galaxyfrom Cowie et al 1983

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Evolution of gas content

Through the Damped Lyα systems,the evolution of the gas in galaxiescould be traced

DLA traced by MgII,between z=0 and 1.6543 systems

n(DLA)(z) ~(1+z)2.27

Rao et al. 1995

Strong evolution, howeverdepends strongly on a fewrare systems at t=0

square = Ω lum, circle =Ω(HI)

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Conclusions

After recombination, GMCs of 10 5-6 Mo collapse and fragmentdown to 10-3 Mo, efficient H2 cooling

The bulk of the gas might not form starsbut a fractal structure, in statistical equilibrium with TCMBLow level, sporadic star formation

==> after first stars, Reheating and reionization

The cold gas survives and will be assembled in larger scalestructures to form galaxies

A way to solve the cooling catastrophy (Blanchard et al 92)

Regulate the consumption of gas into stars (reservoir)