AstrochemistrySpecial course in astronomy 53855, 5 op

Jorma Harju, Julien Montillaud, Olli Sipilä

Department of Physics

Spring 2013, Fridays 12:15-13:45, Lecture room D117Course web page

http://www.courses.physics.helsinki.fi/astro/astrokemia

AstrochemistrySpecial course in astronomy 53855, 5 op

Jorma Harju, Julien Montillaud, Olli Sipilä

Department of Physics

Spring 2013, Fridays 12:15-13:45, Lecture room D117Course web page

http://www.courses.physics.helsinki.fi/astro/astrokemia

OutlineCourse plan

Introduction

Cosmic abundances

Big Bang nucleosynthesisDeuteriumHeliumLithium ja BerylliumPrimordial abundancesRecombination

Nucleosynthesis in StarsLife cycle of a StarLow-mass starsIntermediate-mass starsMassive starsSupernovae

Summary

Course plan

date lecturer topic18.1. Jorma Cosmic abundances and their origin25.1. Jorma Interstellar chemistry, formation of simple molecules1.2. Julien Formation of H2

8.2. Julien Polycyclic Aromatic Hydrocarbons in space15.2. Julien Gas-grain interaction in the interstellar medium22.2. Julien Laboratory astrochemistry

1.3. Jorma Observational astrochemistry8.3. Teaching break

15.3. Jorma Circulation of interstellar matter, molecular havens in space22.3. Jorma Chemistry in cold and hot cores, circumstellar disks

and the formation of prebiotic molecules29.3. Good Friday

5.4. Jorma Primordial chemistry12.4. Olli Astrochemical modelling 119.4. Olli Astrochemical modelling 226.4. Olli Modelling exercise

3.5. Exam

Astrochemistry

I Aims to understand chemical abundances and processes ininterstellar, circumstellar, and planetary environments

I Started in early 1970’s, after the discovery of some simplemolecules in space using radio spectroscopyHollenbach & Salpeter: formation of H2 on dust grains

Klemperer & Herbst: Ion-molecule reactions possible in cold interstellargas, in particular H+

2 + H2 → H+3 + H

I Astrochemical research:-observational studies (abundances in space)-laboratory work (reaction rate coefficients, branching ratios),-theoretical work and modelling (reaction systems)

I A list of intestellar and circumstellar species detected so farhttp://www.astrochymist.org/astrochymist_ism.html

The importance of astrochemistry

I Needed to explain the formation of stars:-molecular line cooling allows the gravitational collapse of gasclouds-regulates the coupling of gas to the magnetic field

I The most abundant molecule in the Universe, H2, is difficult todetect: substitutes with known abundances (tracer molecules)needed

I The astrophysical interpretation of molecular line data is oftenhalting without knowledge of astrochemistry (e.g., degree ofionization, shock properties, isotopic ratios)

I Astrochemistry helps to identify probes for special conditions(e.g., dense nuclei of pre-stellar cores, hot cores around protostars)

I Astrochemistry drives and complements laboratory studies(extreme conditions)

The 10 most abundant elements

Solar system and the local interstellar mediumsolar system local ISM∗

element Z [X]/[H] mass fraction gas, [X]/[H] dust, [X]/[H]1H hydrogen 1 1 0.714He helium 2 0.098 0.28

16O oxygen 8 4.9 10−4 5.6 10−3 2.8 10−4 2.6 10−4

12C carbon 6 2.5 10−4 2.1 10−3 1.8 10−4 2.1 10−4

20Ne neon 10 1.0 10−4 1.4 10−3

14N nitrogen 7 8.5 10−5 8.5 10−4 5.0 10−5 3.6 10−5

28Si silicon 14 3.5 10−5 6.9 10−4 5.0 10−6 2.9 10−5

24Mg magnesium 12 3.5 10−5 5.9 10−4 2.9 10−6 3.2 10−5

56Fe iron 26 2.8 10−5 1.1 10−3 1.4 10−6 2.7 10−5

32S 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 meteoriticsamples (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 modelsfor different spectral classes.Synthetic spectra are compared with observed ones.

Solar system abundances

The abundances for elements heavier than oxygen are similarin the solar photosphere and in some carbon meteorites.These are believed to correspond the situation in the pre-solarnebula, 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 heliumThe abundances of 30-40 elements heavier than He are determinedin the solar neighbourhood (local ISM, within ∼ 1.5 kpc)

Interstellar medium 2

Elemental abundances in theISM relative to those in thesolar system as functions ofthe condensation tempera-ture.Derived from absorption line observations to-

wards ζ Oph(Savage & Sembach 1996, ARA&A 34, 279)

The abundancies of volatile species are within a factor of two the same as inthe 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 thetemperature where dust grains were formed in the stellar wind or supernovaexplosion.Difficulty: the elements are distributed between gases and solid particles

Isotopic ratios

Isotopic ratio Solar system Local ISMD/Hd 3.4 10−5 1.6 10−5

4He/H 9.8 10−2 8.9 10−2

3He/4He 1.4 10−4 ...7Li/Ha 1.9 10−9 ...7Li/6Li 12.3 6-1312C/13C 89b 60-8014N/15N 270b,c 430-47016O/18O 490b 530-59018O/17O 5.5b 3-432S/34S 22b ∼ 22

reference: Wilson & Rood 1994, ARA&A, 32, 191∗ IGM: D/H∼ 2.8 10−5 (primordial)a Population II stars: 7Li/H ∼ 1.7 10−10 (primordial?)b Anders & Grevesse 1989, Geochim. Cosmochim. Acta 53, 197c 250-500 Marti & Kerridge 2010, Science 328, 1112Galactic gradient: 12C/13C, 14N/15N, 16O/18O increase with galactocentric distance

“Periodic table of elements” in astronomy

This is the current situation - in the beginning heavy elementsand dust were missingHow did we end up in this situation?

Origin of elements

General principle

I Z=1-5: Big Bang-All hydrogen (H,D), almost all helium (3He, 4He), and part oflithium (7Li)-the elements with Z=2-4 are generated in stars, H and Ddiminish

I Z=5-26: Fusion reactions in starsI Z=27-94: Supernova explosions

-neutron capture or beta decay-natural 93 Np ja 94 Pu are extremely rare

I Spallation (fission) of interstellar CNO nuclei caused bycosmic particles (p, α) generate most of the elements withZ=4-5

Big Bang nucleosynthesis 1: protons and neutrons

I protons and neutrons (and their antiparticles) were formedin the expanding and cooling Universe about t ∼ 1µs afterBig Bang, when particles in thermal equilibrium had akinetic energy of 1 GeV (T ∼ 1013 K)

I Most of these were annihililated: p + p→ γ + γ,n + n→ γ + γ, but a small amount of nucleons survived

I At high temperatures p and n can convert to each other,e.g., e− + p→ n + νe (endothermic by ∆E = −0.83 MeV).In the cooling Universe β decay started to reduce theabundance of neutrons: n→ e− + p + νe

Big Bang nucleosynthesis 2: deuteriumI 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, 3 109 K)

I At the temperature T ∼ 0.06− 0.07 MeV (7 108 K), there wasenough deuterium for the formation of helium nuclei, 3He++,4He++

Big Bang nucleosynthesis (3): helium

Deuterium reactions (charge signs omitted)d + p→3 He + γ, d + d→3 He + nn +3 He→4 He + γ, d +3 He→4 He + pTritium (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

I In addition to the d, 3He, and 4He nuclei, very smallamounts of 7Li- ja 7Be nuclei were formed:4He +3 He→7 Be + γ4He + t→7 Li + γ7Be + n→7 Li + p

I Lithium can fragment into helium:7Li + p→4 He +4 He

I The unstable t converted to 3He through β decay7Be converted to 7Li through proton capture(both 7Be and 8Be are unstable, stable isotope 9Be)

I 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

I The efficiency of fusion reaction decreases strongly with thetemperature. The primordial fusion was practically stopped atthe formation of He

I When n was bound to d and He nuclei the n/p ratio was frozento ∼ 1/7 (the lifetime of a free n is about 15 min)

I The n/p ratio determines also the 4He/1H ratio ∼ 1/12 (themass fraction of 4He: Y ≈ 0.25)

I The abundance ratios 4He/H, 3He/H, D/H, and 7Li/H dependstrongly on the ratio of baryons and photons (η).(The observed ratios have been used to derive the baryonic densityparameter ΩB, as the photon density can be measure from the cosmicmicrowave background.)

Recombination

The nuclei and electrons combineto form neutral atoms- “recombination”The Universe became transparent:CMB

He was the first to recombine:He++ + e− → He+ + hν (z ∼ 6000, T ∼ 16000 K, t ∼ 20000 yr)

He+ + e− → He + hν (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

I Nuclear fusion requires a very high temperature, andthereby depends on the stellar mass.In brown dwarfs with M ≤ 0.08 M - the “hydrogenburning” cannot start.

I In low-mass stars (M ≤ 1 M, T ≤ 1.5 107 K) hydrogennuclei are converted to helium nuclei in proton-proton chain

I

Main branch (ppI): p + p →d + e+ + νed + p→3 He + γ3He +3 He→4 He + p + p

Nuclear fusion in low-mass stars 2

I Helium is ignited during the giant phase (if M > 0.26M),when the triple-alpha (3α) reaction starts in thedegenerated helium core (Helium flash)

I

A red giant will become awhite dwarf, outer parts areexpelled and dissolve in thesurroundings - planetary neb-ula

Nuclear synthesis in intermediate-mass stars 1

I Intermediate-mass and massive stars convert hydrogen tohelium through the carbon cycle (if C from the ISM areavailable)

C acts here as a catalyst

Nuclear synthesis in intermediate-mass stars 2

I The burning of helium to carbon (3α→12 C) startsgradually during the giant phase, and proceeds from thecore to a burning shell

I The reaction can continue 12C(α, γ)16O, 16O(α, γ)20Ne,especially in massive stars

I In low-mass and intermediate-mass stars (1− 10M) thecarbon core does not become hot enough to be ignited

Nuclear synthesis in massive stars 1

I Carbon is efficiently converted to oxygen in alpha captureI When helium is consumed carbon starts to burn: 12C +12 C.

This produces mainly 20Ne nuclei: 12C(12C, α)20Nebut also 12C(12C, γ)24Mg and 12C(12C, p)23Na

I The burning of neon starts with breaking into oxygen caused byabsorbed gamma-rays (photons)20Ne(γ, α)16OThe α particles (4He nuclei) are recycled:20Ne(α, γ)24Mg , 24Mg(α, γ)28SiSide products: e.g. 27Al, 31P, and 32S

I The principal products of oxygen burning, 16O+16O, are the socalled α nuclei, 28Si, 32S, 36Ar, and 40Ca

Nuclear synthesis in massive stars 2

I The burning of silicon, 28Si, begins (like neon burning) withdisruption induced by a photon.Light nuclei form heavier nuclei nuclei as long as the bindingenergy per nucleon, Q, increases with the massQ = [Zmp + Nmn −m(Z ,N)]c2/AElectrostatic repulsion dampens fusion with increasing ZThe end products are nickel and iron, in short28Si +28 Si→56 Ni + γ, 56Ni→56 Fe + 2e+ + 2νe,

I

The binding energy per nucleonreaches the maximum at 56Fe(cannot yield energy by fusion or fis-sion)

Nuclear reactions in a massive star (M = 20M)

fuel product side products T (109 K) duration (yr) main reactionH He 14N 0.037 8.1 106 4 1H→4 He (CNO cycle)He O, C 18O, 22Ne 0.19 1.2 106 3 4He→12 C

(s-process) 12C(α, γ)16OC Ne, Mg Na 0.87 9.8 102 12C +12 C→ ...

Ne O, Mg Al, P 1.6 0.60 20Ne→16 O +4 He20Ne +4 He→24 Mg

O Si, S Cl, Ar, 2.0 1.3 16O +16 O→ ...K, Ca

Si Fe Ti, V, Cr, 3.3 0.031 28Si→24 Mg +4 He...Mn, Co, Ni 28Si +4 He→24 Mg...

Species heavier than iron can from through the so called slowneutron capture, s-process. This requires iron from the ISM.

Nuclear synthesis in supernovae 1

IA single massive star ends itslife as type II supernova

I The collapse of the Fe nucleus is followed by an expandingshock wave.This starts a series of explosive nuclear reactions, begining withthe breaking up of nuclei into α particles and nucleons.

I α reaktion produce quickly multiples of 4He nuclei up to 64Ge.After 40Ca these are unstable (the isotopes 44Ti, 52Fe, 56Ni, 60Zn,64Ge).After the explosion radioactive decay produces stable isotopes48Ti, 52Cr, 56Fe, ...

Nuclear synthesis in supernovae 2

I A neutron star formed in the centre emits an immense amount ofneutrinos which interact with nuclei.Rare metals are synthesized in reactions where neutrinosconvert a neutron to a protons (νe + n→ p + e−), or remove anucleon:138Ba→138 La , 180Hf→180 Ta ,12C→11 B , 20Ne→19 F

I The heaviest species, A ≥ 130− 140, are believed to result fromthe r-process (rapid neutron capture) in the expanding shellheated by the neutrino flux.

Nuclear synthesis in supernovae 3

I SNe II produce approximately solar system abundances, exceptthat the nuclei 16O - 40Ca are overabundant by a factor of 2-3relative to the range 48Ti - 64Zn.

I Judging from this SNe II produce ∼ 1/3− 1/2 of the elements ofthe "iron peak" (Ti,V,Cr,Mn,Fe,Co,Ni,Cu,Zn)

I The rest is likely to come from type Ia supernovae (explosion ofa white dwarf in a binary).The luminosity peak in SNe Ia is probably caused by the decayof nickel to iron: 56Ni→56 Co→56 Fe(SNe Ia are the brightest stars in the Universe, and important for thedetermining the cosmological distance scale.)

Accounting of cosmic abundances

I H and He most abundant:-formed in Big Bang

I Exponential decrease of abundances with Z:-coulomb repulsion and decreasing gain in α captures

I Elements with even Z more abundant:-α captures

I Low abundances of heavy elements:-slow neutron capture in normal stars

I Low abundances of Li, Be, B:-bypassed by stellar nucleosynthesis-destroyed by nucleon bombardment in stars-replenished by cosmic ray spallation