Astrochemistry Lecture 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 and 8 unstable) -all hydrogen (H,D), major part of helium (3He, 4He), part of lithium (7Li) -more of the elements Z=2-4 are synthesized in stars, while H and D decrease

I Primordial relative abundances: n(4He)/n(H)= 0.083 n(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 have been larger in some regions of the universe, resulting in small amounts 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 − 1.9 Myr (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), cooling adiabatically because of the expansion

I TCBR = 2.7(1 + z) K, ∼ 4000 K at the time of recombination z = 1100, t = 370000 yr −6 2 3 −3 I nH ∼ 8 10 Ωbh (1 + z) cm ∼ 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 possible

I Chemistry is, however, complicated by a large number of possible quantum states owing to collisions and interaction with the CBR (Coppola et a. 2011, ApJS 193, 7) − −4 I The recombination was not complete, X(e ) ∼ 10 -otherwise the chemistry could not have started

I Once neutral He increased in abundance charge transfer with ions became possible

I The first molecules started to form at z ∼ 2000 (t ∼ 10000 yr)

I The abundances “freeze out” by z ∼ 100 owing to + + expansion, with small amounts of H2, HD, H2 , HeH , etc. The first molecules: Helium chemistry

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

+ + He + He → He2 + hν H+ + He → HeH+ + hν + I When H2 ions and H2 molecules are available (see below), also the following reactions can produce HeH+: + + + + H2 + He ↔ HeH + H , H2 + He ↔ HeH + H + I H2 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 + hν is very inefficient because the system does not have time to emit a photon before H2 dissociates

I Principal production pathways of H2 in the early universe: 1) Radiative association and charge transfer: + + H + H → H2 + hν , + + H2 + H → H2 + H + -the latter reaction is fast, the H2 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 through + − photodissociation of H2 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 to Lepp, Stancil & Dalgarno (1998, MmSAI 69, 331) The first molecules: Deuterium chemistry

I 1) Radiative association and charge transfer

D+ + H → HD+ + hν HD+ + H → HD + H+

I 2) Charge transfer and deuteration of H2

D + H+ → D+ + H + + D + 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 transfer LiH+ + 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 LiH 4.6 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 the gas 3 Matter density ρM ∝ (1 + z) The temperature is determined by 1) interaction between the CMB and the electrons, 2) adiabatic expansion of the universe, 3) collisional excitation of atoms and molecules, followed by 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 abundances by 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 according to D. Puy: −6 −9 −19 H2/H ∼ 10 , HD/H ∼ 10 , LiH/H ∼ 10 -the most important coolants in the early universe − −4 I fractional ionization X(e ) ∼ 3 10 . + −12 + −18 Cations: X(H2 ) ∼ 1.3 10 , X(HD ) ∼ 2.1 10 , + −14 + −13 X(H2D ) ∼ 5.1 10 , X(HeH ) ∼ 6.2 10 , X(LiH+) ∼ 9.4 10−18 The first structures

I BB cosmology: the imprints of tiny inhomogeneities in matter density remaining (but enlarged in size) during the cosmic inflation are seen as tempetature variations in the CMB I According to recent Planck results (2013), 68.3% of the universe consist of dark energy, and 31.5% of matter. The share 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 of dark matter has relatively low-mass peaks, “mini-halos”, 6 M ∼ 10 M (figure: Bromm et al. 2009)

I The first, Population III.1 stars were formed from visible matter 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 not been 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 the gas 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 one mini-halo

I The combined effect of Pop III.1 stars and/or active galactic nuclei (AGNs) postponed the formation of the next stellar population (Pop III.2) by ∼ 100 Myr The universe was completely reionized at z ∼ 11 (t ∼ 410000 yr) Simulation: bubbles ionized by the first stars (blue), molecular regions marked with green Ion-molecule chemistry restarted in cooling ionized regions (Bromm et al. 2009) The first stars (4)

I A large fraction of Pop III.1 stars probably collapsed to black 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), sprinkling heavier 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 not possible if no heat is removed: adiabatic collapse results in the ionization of hydrogen, the radiation pressure dissolves the cloud

I Compression of the gas raises the kinetic temperature Gravity must win the gas pressure 3/2 −1/2 Jeans’ mass: MJ ∼ T ρ , T temperature, ρ mass density

I The collapse can continue only if the gas can cool

I The gas can remove thermal energy via radiation

I The Jeans’ mass decreases because of cooling → less massive condensations can collapse -clouds can fragment into smaller, gravitationally bound parts Collapse of a molecular cloud (2)

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

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

I Dipole transitions of HD and LiH were important at T ∼ 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 dissociated in 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 -> radiation

I The radiative cooling of the gas below 1000 K needs molecules

I 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 prerequisite to the formation of the first stars

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

I Massive stars re-ionized the gas, and enriched the ISM through supernova explosions Summary (2)

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

I Molecular clouds are mainly H2 gas. The thermal properties of H2 affect the collapse of molecular clouds to stars

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 emission I The significance of dust: 1) Attenuation of starlight 2) Strong source of radiation at infrared and submillimeter wavelengths 3) Catalyst in H2 production 4) Accretes icy mantles where complex organic molecles can form I Because of the efficient cooling the present-day stellar mass distribution is dominated by low-mass stars