Nucleosynthesis. Lecture 8: Supernovae

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1 Nucleosynthesis Lecture 8: Supernovae

2 Overview Lecture 1: Introduction & overview Lecture 2: Thermonuclear reactions Lecture 3: Big-bang nucleosynthesis Lecture 4: Thermonuclear reactions inside stars I (H-burning) Lecture 5: Thermonuclear reactions inside stars II (advanced burning) Lecture 6: Neutron-capture and supernovae I Lecture 7: Neutron-capture and supernovae II Lecture 8: Thermonuclear supernovae Lecture 9: Li, Be and B Lecture 10: Galactic chemical evolution and relation to astrobiology Paper presentations I June 21 Paper presentations II June 28

3 nucleosynthesis in neutron star mergers Potential sources of ejecta in neutron star mergers tidal ejecta collision heated ejecta ν- and B-driven winds post-merger accretion disk

4 nucleosynthesis in neutron star mergers Primary r-process elements are most likely created in NS merger -like collapsar accretion disks

5 nucleosynthesis in supernovae Stellar explosions are the most important source of enrichment in galaxies

6 Supernovae SN 1054 Yu et al. 2015

7 Supernovae SN 1054 SN 1572

8 supernovae Light echoes from SN 1572 Light echoes from SN 1987A Krause et al. 2008, Rest et al. 2005

9 stellar explosions The simplest model for a supernova explosion is that of a conventional bomb Large amount of energy on the verge of being released + trigger Bombs often fail because they are not triggered in the right way. So do most of our model supernova explosions!!! Still not in a position to explain stellar explosions from first principles

10 microphysics

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nuclear statistical equilibrium A system at any temperature and density will come into equilibrium, provided it is maintained for long enough The nuclear statistical equilibrium most relevant for astrophysics is one in which strong and EM interactions participate. The abundance of each element is given in terms of its nuclear properties and the local plasma conditions X i (A i,z i,t, ) = A 3/2 apple 2 ktm(ai,z i ) µ(ai,z i )+B(A i,z i ) G(A, Z, T ) N A h 2 exp kt µ(a i,z i )=Z i µ p + N i µ n = Z i µ p +(A i Z i )µ n Mass (baryon number is conserved). Because weak interactions do not participate in the equilibrium, charge is also conserved (number of protons, neutrons) X i X i = 1; Y e = X Z i A i X i With the above, each triplet (T,,Y e ) (X <latexit sha1_base64="elfe8ikalqywagpdk0gh6iqll1s=">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</latexit> i,µ p,µ n ) yields a unique solution for Blackboard

12 nuclear statistical equilibrium Credit: ILIADIS

13 nuclear statistical equilibrium Credit: ILIADIS

14 nuclear statistical equilibrium Credit: F. Timmes

15 SNe Ia H- and He- deficient, strong silicon lines Lsun M56Ni > 0.6 Msun vej > 10 3 km/s Ekin ~ erg in both elipticals+spiral rate correlates with star formation in spiral similar peak L and early spectra SN 1994D; credit Nasa/HST

16 SNe Ia Standardisable candles with high luminosities: can be used to measure cosmological distances (Phillips 1995)

17 SNe Ia models The origin of SNe Ia is still debated! Most models include a CO white dwarf that somehow approaches the Chandrasekhar limit Problem: stellar evolution makes CO white dwarfs with masses up to ~1.1 Msun Somehow the white dwarf needs to gain ~ Msun Three chief possibilities single-degenerate model Chandrasekhar mass model Mass accretion from a non-degenerate stellar companion until Chandrasekhar limit is reached. Carbon ignites in the centre due to compressional heating single-degenerate model sub-chandrasekhar mass model Mass accretion from a non-degenerate stellar companion builds a thick He layer on top of CO core. Helium detonation leads to compression and carbon ignition double degenerate model Mass accretion in a merger Each model has its own pros and cons. Still many aspects, including ignition and flame propagation are highly uncertain

18 SNe Ia: SD scenario Accretion from a companion if rate too high: common envelope if rate too low: nova eruptions (mass likely decreases) only seems possible for a small range accretion rates 10-7 Msun/yr for billions of years possibly visible as soft x-ray sources

19 SNe Ia: SD scenario As the mass of the star increases due to accretion, so does the density/temperature in the core 12C+ 12 C rate increases (highly screened!), at some point the energy generation rate becomes larger than the neutrino cooling rate > ignition Initially the extra energy drives convection which extends throughout the entire star > carbon shimmering As temperature increases, there comes a point at which τburn ~ τconv Runaway ensues! Burning becomes faster than the time it takes a sound wave to cross a pressure scale hight

20 SNe Ia: SD scenario Ignition may start off-centre and propagate inwards

21 SNe Ia: deflagration, detonation or both?

22 sub-chandrasekhar mass models Helium layer should be thin enough to be invisible in the spectrum but thick enough to ignite ( Msol) Diversity of possible outcomes Can have multiple ignition points

23 mergers Wide range of possible outcomes ranging from prompt detonation to a stable configuration depending on initial masses and compositions

24 SNe Ia nucleosynthesis Consider simple detonation model and post-shock adiabatic cooling Credit: ILIADIS

25 SNe Ia nucleosynthesis Credit: ILIADIS

26 Pre-supernova structure and core-collapse Credit: ILIADIS

27 Pre-supernova structure and core-collapse Credit: ILIADIS

28 Core-collapse supernovae nucleosynthesis neutronization burst Explosion phase Breakout Ejecta accretion shock nucleosynthesis shock revival Fallback Neutrino sphere Proto-NS Bethe & Wilson (1985)