AstrochemistryLecture 10, Primordial chemistry

Jorma Harju

Department of Physics

Friday, April 5, 2013, 12:15-13:45, Lecture room D117

The first atoms (1)

I SBBN (Standard Big Bang Nucleosynthesis): elements Z=1-4(and negligible amounts of heavier elements, the nuclei with A = 5 and8 unstable)

-all hydrogen (H,D), major part of helium (3He, 4He), part oflithium (7Li)-more of the elements Z=2-4 are synthesized in stars, while Hand D decrease

I Primordial relative abundances:n(4He)/n(H)= 0.083n(D)/n(H)= 2.7 10−5,n(7Li)/n(H)= 1.7 10−10,n(3He)/n(H)∼ 0.3 10−5

I Non-standard cosmology: the baryon-to-photon ratio can havebeen larger in some regions of the universe, resulting in smallamounts of C, N, O, and F nuclei

The first atoms (2)

Nuclei recombined with electrons in the cooling universe

I He++ → He+ → He: z ∼ 6000→ 2700,T ∼ 20000→ 10000 K, t ∼ 18000− 78000 yr(IHe = 24.6 eV)

I H+ → H: z ∼ 1100, T ∼ 4000 K , t ∼ 370000 yr(IH = 13.6 eV)

I Li+ → Li z ∼ 500− 400, T ∼ 1900− 1500 K, t ∼ 1.4− 1.9Myr(ILi = 5.4 eV)

Conditions in the early universe

I Gas composition after recombination: H, He, traces of e−, D,and Li

I Gas exposed to the cosmic background radiation (CBR), coolingadiabatically because of the expansion

I TCBR = 2.7(1 + z) K,∼ 4000 K at the time of recombination z = 1100, t = 370000 yr

I nH ∼ 8 10−6Ωbh2(1 + z)3 cm−3

∼ 200 cm−3 at z ∼ 1100

I The gas temperature after recombination, i.e. up to z ∼ 1100:T ∝ (1 + z)2 (Lepp et al. 1998)

Chemistry in the early universe

I Only few reactions possibleI Chemistry is, however, complicated by a large number of

possible quantum states owing to collisions and interactionwith the CBR (Coppola et a. 2011, ApJS 193, 7)

I The recombination was not complete, X (e−) ∼ 10−4

-otherwise the chemistry could not have startedI Once neutral He increased in abundance charge transfer

with ions became possibleI The first molecules started to form at z ∼ 2000 (t ∼ 10000 yr)

I The abundances “freeze out” by z ∼ 100 owing toexpansion, with small amounts of H2, HD, H+

2 , HeH+, etc.

The first molecules: Helium chemistry

I The first molecules are likely to be helium compounds, forexample (radiative association):

He+ + He → He+2 + hνH+ + He → HeH+ + hν

I When H+2 ions and H2 molecules are available (see below),

also the following reactions can produce HeH+:H+

2 + He↔ HeH+ + H , H2 + He+ ↔ HeH+ + HI H+

2 ions react, however, preferentially with H to form H2

The first molecules: Hydrogen chemistry (1)

I A radiative association between two H atoms,H + H→ H2 + hνis very inefficient because the system does not have timeto emit a photon before H2 dissociates

I Principal production pathways of H2 in the early universe:1) Radiative association and charge transfer:H+ + H→ H+

2 + hν ,H+

2 + H→ H2 + H+

-the latter reaction is fast, the H+2 abundance remains low

The first molecules: Hydrogen chemistry (2)

I 2) Catalytic electron attachment:H + e− → H− + hν ,H− + H→ H2 + e−

I The CBR limits the productivity of these reactions throughphotodissociation of H+

2 and photodetachment of H−

(radiation energy density ∝ (1 + z)4)I H2 is mainly dissociated in collisions

The first molecules: Hydrogen chemistry (3)

Hydrogen chemistry network in the early universe according toLepp, Stancil & Dalgarno (1998, MmSAI 69, 331)

The first molecules: Deuterium chemistry

I 1) Radiative association and charge transfer

D+ + H → HD+ + hνHD+ + H → HD + H+

I 2) Charge transfer and deuteration of H2

D + H+ → D+ + HD+ + H2 → HD + H+

I The second pathway is dominant when there is enough H2- main route to HD in diffuse interstellar clouds

The first molecules: Lithium chemistry

I Radiative association 1)Li+ + H→ LiH+ + hν

I or 2)Li + H+ → LiH+ + hν

I Charge transferLiH+ + H→ LiH + H+

-faster than radiative association between Li and H

The first molecules

Fractional abundances at z = 10 according to Puy & Pfenniger (2006):HeH+ H2 HD LiH4.6 10−14 1.13 10−6 3.67 10−10 2.53 10−20

The evolution of molecular abundances (1)

I Chemistry is affected by the density and temperature of thegasMatter density ρM ∝ (1 + z)3

The temperature is determined by1) interaction between the CMB and the electrons,2) adiabatic expansion of the universe,3) collisional excitation of atoms and molecules, followedby radiation,4) heat produced or absorbed by chemical reactions

I The expansion causes that molecular abundances “freeze”at an early stage

I Primordial molecules reach their equilibrium abundancesby the redshift z ∼ 100 (t ∼ 17.5 Myr, D. Puy)

The evolution of molecular abundances (2)

I The abundances of H2, HD, and LiH at z = 100 accordingto D. Puy:H2/H ∼ 10−6, HD/H ∼ 10−9, LiH/H ∼ 10−19

-the most important coolants in the early universeI fractional ionization X (e−) ∼ 3 10−4.

Cations: X (H+2 ) ∼ 1.3 10−12, X (HD+) ∼ 2.1 10−18,

X (H2D+) ∼ 5.1 10−14, X (HeH+) ∼ 6.2 10−13,X (LiH+) ∼ 9.4 10−18

The first structuresI BB cosmology: the imprints of tiny inhomogeneities in

matter density remaining (but enlarged in size) during thecosmic inflation are seen as tempetature variations in theCMB

I According to recent Planck results (2013), 68.3% of theuniverse consist of dark energy, and 31.5% of matter. Theshare of ordinary matter is 4.9% (15.5% of all matter).

The first stars (1)

I Dark matter determines where the detectable (baryonic)matter concentrates

I Cold Dark Matter (CDM) model: the density distribution ofdark matter has relatively low-mass peaks, “mini-halos”,M ∼ 106M (figure: Bromm et al. 2009)

I The first, Population III.1 stars were formed from visiblematter accreted to mini-halos at z ∼ 20− 30(t ∼ 100− 200 Myr).

The first stars (2)

I Pop III stars is a theoretical prediction, they haved notbeen observed (figure: Bromm et al. 2009)

I Pop III stars formed before the first galaxies

The first stars (3)

I Pop III.1 stars were massive, perhaps 60-300 M.-An enormously intensive radiation field could ionize thegas up to radii of several kpc, and photodissociate H2(through Lyman and Werner bands) in much larger regions-Probably only one or a few Pop III.1 were formed in onemini-halo

I The combined effect of Pop III.1 stars and/or activegalactic nuclei (AGNs) postponed the formation of the nextstellar population (Pop III.2) by ∼ 100 MyrThe universe was completely reionized at z ∼ 11(t ∼ 410000 yr)

Simulation: bubbles ionized by thefirst stars (blue), molecular regionsmarked with greenIon-molecule chemistry restarted incooling ionized regions (Bromm etal. 2009)

The first stars (4)

I A large fraction of Pop III.1 stars probably collapsed toblack holes

I In the mass range M ∼ 140− 260M they, however,exploded as supernovae (PISN, pair-instability supernova -gamma radiation is converted to positron-electron pairs in the nucleus,the nucleus collapses, and the whole star explodes), sprinklingheavier elements into space

I Population II stars (majority of stars in globular clusters)have low metallicities

Collapse of a molecular cloud (1)

I The gravitational collapse of density enhancements is notpossible if no heat is removed: adiabatic collapse results inthe ionization of hydrogen, the radiation pressure dissolvesthe cloud

I Compression of the gas raises the kinetic temperatureGravity must win the gas pressureJeans’ mass: MJ ∼ T 3/2ρ−1/2, T temperature, ρ massdensity

I The collapse can continue only if the gas can coolI The gas can remove thermal energy via radiationI The Jeans’ mass decreases because of cooling→ less massive condensations can collapse-clouds can fragment into smaller, gravitationally boundparts

Collapse of a molecular cloud (2)

I According to simulations the first molecular clouds werecreated at the redshift z ∼ 20 (180 Myr)

I In spite of its low abundance, H2 was an efficient coolant attemperatures of T ∼ 200− 10000 K

I Dipole transitions of HD and LiH were important atT ∼ 100− 200 K

I At large densities three-body reactions become feasible:H + H + H→ H2 + H or H2 + H + H→ H2 + H2

I On the other hand, at very high densities H2 is dissociatedin collisions: H2 + H + H→ 4H

Simulations: Stacy et al. (2010)

Simulations: Stacy et al. (2010)

Stellar mass distribution

Present day population: most stars have low masses(0.1M < M < 1M)

Cooling by radiation

I Electronic states of atoms and molecules E ∼ 1 eV(T ∼ 10000K)Vibrational states of molecules: E ∼ 0.1 eV (T ∼ 1000K)Rotational states of molecules: E ∼ 0.01 eV (T ∼ 100K)

I Cooling: collision -> excitation -> radiationI The radiative cooling of the gas below 1000 K needs

moleculesI At large densities and low temperatures (below 10 K)

thermal dust emission is important for cooling

Summary (1)

I The formation of hydrogen molecule, H2, was prerequisiteto the formation of the first stars

I The symmetry of H2 and the implied weak interaction withradiation kept the cooling inefficient.Therefore the first stars were very massive

I Massive stars re-ionized the gas, and enriched the ISMthrough supernova explosions

Summary (2)

I Heavy elements synthesized in stars, and dust formed inthe circumstellar envelopes created preconditions for theformation cool and dense molecular clouds

I Molecular clouds are mainly H2 gas. The thermalproperties of H2 affect the collapse of molecular clouds tostars

I The ionization of H2 initiates interstellar chemistry

Summary (3)I The most important coolants in molecular clouds: CO, C,

O, and O2, thermal dust emissionI The significance of dust:

1) Attenuation of starlight2) Strong source of radiation at infrared and submillimeterwavelengths3) Catalyst in H2 production4) Accretes icy mantles where complex organic moleclescan form

I Because of the efficient cooling the present-day stellarmass distribution is dominated by low-mass stars