Astrochemistry. Special course in astronomy 53855, 5 op. Jorma Harju, Julien Montillaud, Olli Sipilä. Department of Physics

Astrochemistry Special course in astronomy 53855, 5 op Jorma Harju, Julien Montillaud, Olli Sipilä Department of Physics Spring 2013, Fridays 12:15-13:45, Lecture room D117 Course web page http://www.courses.physics.helsinki.fi/astro/astrokemia

Outline Course plan Introduction Cosmic abundances Big Bang nucleosynthesis Deuterium Helium Lithium ja Beryllium Primordial abundances Recombination Nucleosynthesis in Stars Life cycle of a Star Low-mass stars Intermediate-mass stars Massive stars Supernovae Summary

Course plan date lecturer topic 18.1. Jorma Cosmic abundances and their origin 25.1. Jorma Interstellar chemistry, formation of simple molecules 1.2. Julien Formation of H 2 8.2. Julien Polycyclic Aromatic Hydrocarbons in space 15.2. Julien Gas-grain interaction in the interstellar medium 22.2. Julien Laboratory astrochemistry 1.3. Jorma Observational astrochemistry 8.3. Teaching break 15.3. Jorma Circulation of interstellar matter, molecular havens in space 22.3. Jorma Chemistry in cold and hot cores, circumstellar disks and the formation of prebiotic molecules 29.3. Good Friday 5.4. Jorma Primordial chemistry 12.4. Olli Astrochemical modelling 1 19.4. Olli Astrochemical modelling 2 26.4. Olli Modelling exercise 3.5. Exam

Astrochemistry Aims to understand chemical abundances and processes in interstellar, circumstellar, and planetary environments Started in early 1970 s, after the discovery of some simple molecules in space using radio spectroscopy Hollenbach & Salpeter: formation of H 2 on dust grains Klemperer & Herbst: Ion-molecule reactions possible in cold interstellar gas, in particular H + 2 + H 2 H + 3 + H Astrochemical research: -observational studies (abundances in space) -laboratory work (reaction rate coefficients, branching ratios), -theoretical work and modelling (reaction systems) A list of intestellar and circumstellar species detected so far http://www.astrochymist.org/astrochymist_ism.html

The importance of astrochemistry Needed to explain the formation of stars: -molecular line cooling allows the gravitational collapse of gas clouds -regulates the coupling of gas to the magnetic field The most abundant molecule in the Universe, H 2, is difficult to detect: substitutes with known abundances (tracer molecules) needed The astrophysical interpretation of molecular line data is often halting without knowledge of astrochemistry (e.g., degree of ionization, shock properties, isotopic ratios) Astrochemistry helps to identify probes for special conditions (e.g., dense nuclei of pre-stellar cores, hot cores around protostars) Astrochemistry drives and complements laboratory studies (extreme conditions)

The 10 most abundant elements Solar system and the local interstellar medium solar system local ISM element Z [X]/[H] mass fraction gas, [X]/[H] dust, [X]/[H] 1 H hydrogen 1 1 0.71 4 He helium 2 0.098 0.28 16 O oxygen 8 4.9 10 4 5.6 10 3 2.8 10 4 2.6 10 4 12 C carbon 6 2.5 10 4 2.1 10 3 1.8 10 4 2.1 10 4 20 Ne neon 10 1.0 10 4 1.4 10 3 14 N nitrogen 7 8.5 10 5 8.5 10 4 5.0 10 5 3.6 10 5 28 Si silicon 14 3.5 10 5 6.9 10 4 5.0 10 6 2.9 10 5 24 Mg magnesium 12 3.5 10 5 5.9 10 4 2.9 10 6 3.2 10 5 56 Fe iron 26 2.8 10 5 1.1 10 3 1.4 10 6 2.7 10 5 32 S sulfur 16 2.1 10 5 4.9 10 4 1.1 10 5 1.0 10 5 Kimura et al. 2003, ApJ 582, 846

Abundance determinations 1 Spectroscopic observations of the Sun and stars (starting from G. Kirchoff & R. Bunsen 1859, H.N. Russell 1929) Laboratory measurements of the Earth minerals and meteoritic samples (W.D. Harkins 1917, V.M Goldsmith 1938,...) Abundances in the Sun and on the Earth are similar (except for: H,He,Li,C,N,O, noble gases)

Abundance determinations 2 Reliable abundance determinations need atmospheric models for different spectral classes. Synthetic spectra are compared with observed ones.

Solar system abundances The abundances for elements heavier than oxygen are similar in the solar photosphere and in some carbon meteorites. These are believed to correspond the situation in the pre-solar nebula, about 4.6 billion years ago. Astronomers chemistry (mass fractions): X=0.735 (H), Y=0.248 (He), Z=0.017 ( metals )

Interstellar medium 1 One of the standard objects: ζ Ophiuchi (O9III, d140 pc, Hubble UV spectrograph) Interstellar gas is mainly composed of hydrogen and helium The abundances of 30-40 elements heavier than He are determined in the solar neighbourhood (local ISM, within 1.5 kpc)

Interstellar medium 2 Elemental abundances in the ISM relative to those in the solar system as functions of the condensation temperature. Derived from absorption line observations towards ζ Oph (Savage & Sembach 1996, ARA&A 34, 279) The abundancies of volatile species are within a factor of two the same as in the Sun. The depletion on the right is caused by condensation into the dust grains. The distibution is likely to correspond to the phase equilibrium at the temperature where dust grains were formed in the stellar wind or supernova explosion. Difficulty: the elements are distributed between gases and solid particles

Isotopic ratios Isotopic ratio Solar system Local ISM D/H d 3.4 10 5 1.6 10 5 4 He/H 9.8 10 2 8.9 10 2 3 He/ 4 He 1.4 10 4... 7 Li/H a 1.9 10 9... 7 Li/ 6 Li 12.3 6-13 12 C/ 13 C 89 b 60-80 14 N/ 15 N 270 b,c 430-470 16 O/ 18 O 490 b 530-590 18 O/ 17 O 5.5 b 3-4 32 S/ 34 S 22 b 22 reference: Wilson & Rood 1994, ARA&A, 32, 191 IGM: D/H 2.8 10 5 (primordial) a Population II stars: 7 Li/H 1.7 10 10 (primordial?) b Anders & Grevesse 1989, Geochim. Cosmochim. Acta 53, 197 c 250-500 Marti & Kerridge 2010, Science 328, 1112 Galactic gradient: 12 C/ 13 C, 14 N/ 15 N, 16 O/ 18 O increase with galactocentric distance

Periodic table of elements in astronomy This is the current situation - in the beginning heavy elements and dust were missing How did we end up in this situation?

Origin of elements General principle Z=1-5: Big Bang -All hydrogen (H,D), almost all helium ( 3 He, 4 He), and part of lithium ( 7 Li) -the elements with Z=2-4 are generated in stars, H and D diminish Z=5-26: Fusion reactions in stars Z=27-94: Supernova explosions -neutron capture or beta decay -natural 93 Np ja 94 Pu are extremely rare Spallation (fission) of interstellar CNO nuclei caused by cosmic particles (p, α) generate most of the elements with Z=4-5

Big Bang nucleosynthesis 1: protons and neutrons protons and neutrons (and their antiparticles) were formed in the expanding and cooling Universe about t 1µs after Big Bang, when particles in thermal equilibrium had a kinetic energy of 1 GeV (T 10 13 K) Most of these were annihililated: p + p γ + γ, n + n γ + γ, but a small amount of nucleons survived At high temperatures p and n can convert to each other, e.g., e + p n + ν e (endothermic by E = 0.83 MeV). In the cooling Universe β decay started to reduce the abundance of neutrons: n e + p + ν e

Big Bang nucleosynthesis 2: deuterium Deuterium nuclei started to form in a fusion reaction p + n d + γ (exothermic E = 2.22 MeV) This reaction became favoured at t 100 s (T < 0.3 MeV, 3 10 9 K) At the temperature T 0.06 0.07 MeV (7 10 8 K), there was enough deuterium for the formation of helium nuclei, 3 He ++, 4 He ++

Big Bang nucleosynthesis (3): helium Deuterium reactions (charge signs omitted) d + p 3 He + γ, d + d 3 He + n n + 3 He 4 He + γ, d + 3 He 4 He + p Tritium (t): n + d t + n, d + d t + p, n + 3 He t + p p + t 4 He + γ, d + t 4 He + n

Big Bang nucleosynthesis 4: lithium and beryllium In addition to the d, 3 He, and 4 He nuclei, very small amounts of 7 Li- ja 7 Be nuclei were formed: 4 He + 3 He 7 Be + γ 4 He + t 7 Li + γ 7 Be + n 7 Li + p Lithium can fragment into helium: 7 Li + p 4 He + 4 He The unstable t converted to 3 He through β decay 7 Be converted to 7 Li through proton capture (both 7 Be and 8 Be are unstable, stable isotope 9 Be) In nature there is no element with the mass number A = 5, and no stable nucleus with the mass number A = 8

Primordial abundances of elements The efficiency of fusion reaction decreases strongly with the temperature. The primordial fusion was practically stopped at the formation of He When n was bound to d and He nuclei the n/p ratio was frozen to 1/7 (the lifetime of a free n is about 15 min) The n/p ratio determines also the 4 He/ 1 H ratio 1/12 (the mass fraction of 4 He: Y 0.25) The abundance ratios 4 He/H, 3 He/H, D/H, and 7 Li/H depend strongly on the ratio of baryons and photons (η). (The observed ratios have been used to derive the baryonic density parameter Ω B, as the photon density can be measure from the cosmic microwave background.)

Recombination The nuclei and electrons combine to form neutral atoms - recombination The Universe became transparent: CMB He was the first to recombine: He ++ + e He + + hν He + + e He + hν (z 6000, T 16000 K, t 20000 yr) (z 2700, T 7000 K, t 80000 yr) H: H + + e H + hν (z 1300, T 3600 K, t 380000 yr) Li: z 500 (T 1400 K, t 1.3 Myr)

Life cycle of a star Nucleosynthesis continues in stars

Nuclear fusion in low-mass stars 1 Nuclear fusion requires a very high temperature, and thereby depends on the stellar mass. In brown dwarfs with M 0.08 M - the hydrogen burning cannot start. In low-mass stars (M 1 M, T 1.5 10 7 K) hydrogen nuclei are converted to helium nuclei in proton-proton chain Main branch (ppi): p + p d + e + + ν e d + p 3 He + γ 3 He + 3 He 4 He + p + p

Nuclear fusion in low-mass stars 2 Helium is ignited during the giant phase (if M > 0.26M ), when the triple-alpha (3α) reaction starts in the degenerated helium core (Helium flash) A red giant will become a white dwarf, outer parts are expelled and dissolve in the surroundings - planetary nebula

Nuclear synthesis in intermediate-mass stars 1 Intermediate-mass and massive stars convert hydrogen to helium through the carbon cycle (if C from the ISM are available) C acts here as a catalyst

Nuclear synthesis in intermediate-mass stars 2 The burning of helium to carbon (3α 12 C) starts gradually during the giant phase, and proceeds from the core to a burning shell The reaction can continue 12 C(α, γ) 16 O, 16 O(α, γ) 20 Ne, especially in massive stars In low-mass and intermediate-mass stars (1 10M ) the carbon core does not become hot enough to be ignited

Nuclear synthesis in massive stars 1 Carbon is efficiently converted to oxygen in alpha capture When helium is consumed carbon starts to burn: 12 C + 12 C. This produces mainly 20 Ne nuclei: 12 C( 12 C, α) 20 Ne but also 12 C( 12 C, γ) 24 Mg and 12 C( 12 C, p) 23 Na The burning of neon starts with breaking into oxygen caused by absorbed gamma-rays (photons) 20 Ne(γ, α) 16 O The α particles ( 4 He nuclei) are recycled: 20 Ne(α, γ) 24 Mg, 24 Mg(α, γ) 28 Si Side products: e.g. 27 Al, 31 P, and 32 S The principal products of oxygen burning, 16 O+ 16 O, are the so called α nuclei, 28 Si, 32 S, 36 Ar, and 40 Ca

Nuclear synthesis in massive stars 2 The burning of silicon, 28 Si, begins (like neon burning) with disruption induced by a photon. Light nuclei form heavier nuclei nuclei as long as the binding energy per nucleon, Q, increases with the mass Q = [Zm p + Nm n m(z, N)]c 2 /A Electrostatic repulsion dampens fusion with increasing Z The end products are nickel and iron, in short 28 Si + 28 Si 56 Ni + γ, 56 Ni 56 Fe + 2e + + 2ν e, The binding energy per nucleon reaches the maximum at 56 Fe (cannot yield energy by fusion or fission)

Nuclear reactions in a massive star (M = 20M ) fuel product side products T (10 9 K) duration (yr) main reaction H He 14 N 0.037 8.1 10 6 4 1H 4 He (CNO cycle) He O, C 18 O, 22 Ne 0.19 1.2 10 6 3 4 He 12 C (s-process) 12 C(α, γ) 16 O C Ne, Mg Na 0.87 9.8 10 2 12 C + 12 C... Ne O, Mg Al, P 1.6 0.60 20 Ne 16 O + 4 He 20 Ne + 4 He 24 Mg O Si, S Cl, Ar, 2.0 1.3 16 O + 16 O... K, Ca Si Fe Ti, V, Cr, 3.3 0.031 28 Si 24 Mg + 4 He... Mn, Co, Ni 28 Si + 4 He 24 Mg... Species heavier than iron can from through the so called slow neutron capture, s-process. This requires iron from the ISM.

Nuclear synthesis in supernovae 1 A single massive star ends its life as type II supernova The collapse of the Fe nucleus is followed by an expanding shock wave. This starts a series of explosive nuclear reactions, begining with the breaking up of nuclei into α particles and nucleons. α reaktion produce quickly multiples of 4 He nuclei up to 64 Ge. After 40 Ca these are unstable (the isotopes 44 Ti, 52 Fe, 56 Ni, 60 Zn, 64 Ge). After the explosion radioactive decay produces stable isotopes 48 Ti, 52 Cr, 56 Fe,...

Nuclear synthesis in supernovae 2 A neutron star formed in the centre emits an immense amount of neutrinos which interact with nuclei. Rare metals are synthesized in reactions where neutrinos convert a neutron to a protons (ν e + n p + e ), or remove a nucleon: 138 Ba 138 La, 180 Hf 180 Ta, 12 C 11 B, 20 Ne 19 F The heaviest species, A 130 140, are believed to result from the r-process (rapid neutron capture) in the expanding shell heated by the neutrino flux.

Nuclear synthesis in supernovae 3 SNe II produce approximately solar system abundances, except that the nuclei 16 O - 40 Ca are overabundant by a factor of 2-3 relative to the range 48 Ti - 64 Zn. Judging from this SNe II produce 1/3 1/2 of the elements of the "iron peak" (Ti,V,Cr,Mn,Fe,Co,Ni,Cu,Zn) The rest is likely to come from type Ia supernovae (explosion of a white dwarf in a binary). The luminosity peak in SNe Ia is probably caused by the decay of nickel to iron: 56 Ni 56 Co 56 Fe (SNe Ia are the brightest stars in the Universe, and important for the determining the cosmological distance scale.)

Accounting of cosmic abundances H and He most abundant: -formed in Big Bang Exponential decrease of abundances with Z: -coulomb repulsion and decreasing gain in α captures Elements with even Z more abundant: -α captures Low abundances of heavy elements: -slow neutron capture in normal stars Low abundances of Li, Be, B: -bypassed by stellar nucleosynthesis -destroyed by nucleon bombardment in stars -replenished by cosmic ray spallation