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Neutrinos and the stars

Neutrinos and the stars. Neutrinos and the Stars. Georg Raffelt, MPI for Physics Lectures at the Topical Seminar Neutrino Physics & Astrophysics 17 - 21 Sept 2008, Beijing, China. Where do Neutrinos Appear in Nature?. Nuclear Reactors. . Sun. . Supernovae (Stellar Collapse).

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Neutrinos and the stars

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  1. Neutrinos and the stars Neutrinos and the Stars Georg Raffelt, MPI for Physics Lectures at the Topical Seminar Neutrino Physics & Astrophysics 17-21 Sept 2008, Beijing, China

  2. Where do Neutrinos Appear in Nature? Nuclear Reactors  Sun  Supernovae (Stellar Collapse) Particle Accelerators  SN 1987A Earth Atmosphere (Cosmic Rays)  Astrophysical Accelerators Soon ? Earth Crust (Natural Radioactivity)  Cosmic Big Bang (Today 330 n/cm3) Indirect Evidence

  3. Where do Neutrinos Appear in Nature? Neutrinos from nuclear reactions: Energies 1-20 MeV Quasi thermal sources Supernova: T ~ few MeV Big-Bang Neutrinos: Very small energies today (cosmic red shift) Like matter today • “Beam dump neutrinos” • High-energy protons hit • matter or photons • Produce secondary p • Neutrinos from pion • decay • p  m + nm • me + nm+ne • Energies ≫ GeV

  4. Where do Neutrinos Appear in Nature? Low-energy neutrino astronomy (including geo-neutrinos) Energies ~ 1-50 MeV • Long-baseline • neutrino oscillation • experiments with • Reactor neutrinos • Neutrino beams from • accelerators • Precision cosmology & • limit on neutrino mass • Big-bangnucleosynthesis • Leptogenesis High-energy neutrino astronomy Closely related to cosmic-ray physics

  5. Neutrinos from the Sun Hans Bethe (1906-2005, Nobel prize 1967) Thermonuclear reaction chains (1938) Helium Reaction- chains Energy 26.7 MeV Solar radiation: 98 % light 2 % neutrinos At Earth 66 billion neutrinos/cm2 sec

  6. Bethe’s Classic Paper on Nuclear Reactions in Stars No neutrinos from nuclear reactions in 1938 …

  7. Gamow & Schoenberg, Phys. Rev. 58:1117 (1940)

  8. Gamow & Schoenberg 2

  9. Sun Glasses for Neutrinos? 8.3 light minutes Several light years of lead needed to shield solar neutrinos Bethe & Peierls 1934: “… this evidently means that one will never be able to observe a neutrino.”

  10. First Detection (1954 -1956) Anti-Electron Neutrinos from Hanford Nuclear Reactor 3 Gammas in coincidence n Cd p e+ e- g g g Clyde Cowan (1919 – 1974) Fred Reines (1918 – 1998) Nobel prize 1995 Detector prototype

  11. First Measurement of Solar Neutrinos Inverse beta decay of chlorine 600 tons of Perchloroethylene Homestake solar neutrino observatory (1967-2002)

  12. Neutrinos from the Sun Solar Neutrinos

  13. Hydrogen burning: Proton-Proton Chains < 0.420 MeV 1.442 MeV 100% 0.24% 85% 15% PP-I hep < 18.8 MeV 90% 10% 0.02% 0.862 MeV 0.384 MeV < 15 MeV PP-II PP-III

  14. Solar Neutrino Spectrum

  15. Hydrogen Burning: CNO Cycle (p,a) (p,a) (p,g) (p,g) (p,g) (p,g) (p,g)

  16. Missing Neutrinos from the Sun Homestake Chlorine 8B Calculation of expected experimental counting rate from various source reactions CNO 7Be Measurement (1970–1995) John Bahcall 1934 - 2005 Raymond Davis Jr. 1914 - 2006

  17. Results of Chlorine Experiment Average Rate Average (1970-1994)2.56  0.16stat 0.16sysSNU (SNU = Solar Neutrino Unit = 1 Absorption / sec / 1036 Atoms) Theoretical Prediction 6-9 SNU “Solar Neutrino Problem” since 1968

  18. Neutrino Flavor Oscillations Two-flavor mixing Each mass eigenstate propagates as with Phase difference implies flavor oscillations Probabilitynenm sin2(2q) Bruno Pontecorvo (1913 – 1993) Invented nu oscillations z Oscillation Length

  19. Cherenkov Effect Light Electron or Muon (Charged Particle) Neutrino Light Cherenkov Ring Elastic scattering or CC reaction Water

  20. Super-Kamiokande Neutrino Detector 42 m 39.3 m

  21. Super-Kamiokande: Sun in the Light of Neutrinos

  22. 2002 Physics Nobel Prize for Neutrino Astronomy Ray Davis Jr. (1914 - 2006) Masatoshi Koshiba (*1926) “for pioneering contributions to astrophysics, in particular for the detection of cosmic neutrinos”

  23. Solar Neutrino Spectrum 7-Be line measured by Borexino (since 2007)

  24. Solar Neutrino Spectroscopy with BOREXINO • Neutrino electron scattering • Liquid scintillator technology • (~ 300 tons) • Low energy threshold • (~ 60 keV) • Online since 16 May 2007 • Expected without flavor oscillations 75 ± 4 counts/100t/d • Expected with oscillations 49 ± 4 counts/100t/d • BOREXINO result (May 2008) 49 ± 3stat ± 4syscnts/100t/d arXiv:0805.3843 (25 May 2008)

  25. Next Steps in Borexino • Collect more statistics of Beryllium line • Seasonal variation of rate • (Earth orbit eccentricity) • Measure neutrinos from the CNO reaction chain • Information about solar metal abundance Measure geo-neutrinos (from natural radioactivity in the Earth crust) Approx. 7-17 events/year Main background: Reactors ~ 20 events/year

  26. Geo Neutrinos: Why and What? • We know surprisingly little about • the interior of the Earth: • Deepest bore hole ~ 12 km • Samples from the crust are • available for chemical analysis • (e.g. vulcanoes) • Seismology reconstructs density • profile throughout the Earth • Heat flow from measured • temperature gradients 30-44 TW • (BSE canonical model, based on • cosmo-chemical arguments, • predicts ~ 19 TW from crust and • mantle, none from core) • Neutrinos escape freely • Carry information about chemical composition, radioactive heat production, • or even a putative natural reactor at the core

  27. Expected Geo Neutrino Fluxes S. Dye, Talk 5/25/2006 Baltimore

  28. Geo Neutrinos Predicted geo neutrino flux KamLAND scintillator detector (1 kton) Reactor background

  29. Kamland Observation of Geoneutrinos • First tentative observation of geoneutrinos • at Kamland in 2005 (~ 2 sigma effect) • Very difficult because of large background • of reactor neutrinos • (is main purpose for neutrino oscillations)

  30. Neutrinos from the Sun Solar Models

  31. Equations of Stellar Structure Hydrostatic equilibrium r P GN r Mr Lr e Radius from center Pressure Newton’s constant Mass density Integrated mass up to r Luminosity (energy flux) Local rate of energy generation [erg/g/s] Energy conservation Energy transfer Opacity Radiative opacity Electron conduction Assume spherical symmetry and static structure (neglect kinetic energy) Excludes: Rotation, convection, magnetic fields, supernova-dynamics, … • Literature • Clayton: Principles of stellar evolution and • nucleosynthesis (Univ. Chicago Press 1968) • Kippenhahn & Weigert: Stellar structure • and evolution (Springer 1990)

  32. Convection in Main-Sequence Stars Sun Kippenhahn & Weigert, Stellar Structure and Evolution

  33. Virial Theorem and Hydrostatic Equilibrium Hydrostatic equilibrium Integrate both sides L.h.s. partial integration with P = 0 at surface R Classical monatomic gas: (U density of internal energy) Average energy of single “atoms” of the gas Virial Theorem Most important tool to understand self-gravitating systems

  34. Virial Theorem Applied to the Sun Virial Theorem Approximate Sun as a homogeneous sphere with Mass Radius Gravitational potential energy of a proton near center of the sphere Thermal velocity distribution Estimated temperature T = 1.1 keV Central temperature from standard solar models

  35. Constructing a Solar Model: Fixed Inputs Solve stellar structure equations with good microphysics, starting from a zero-age main-sequence model (chemically homogeneous star) to present age Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  36. Constructing a Solar Model: Free Parameters • 3 free parameters • Convection theory has 1 free parameter: • Mixing length parameter aMLT • determines the temperature stratification where convection • is not adiabatic (upper layers of solar envelope) • 2 of the 3 quantities determining the initial composition: • Xini, Yini, Zini (linked by Xini + Yini + Zini = 1). • Individual elements grouped in Zini have relative abundances • given by solar abundance measurements (e.g. GS98, AGS05) • Construct a 1 M⊙initial model with Xini, Zini, (Yini = 1 -–Xini - Zini) • and aMLT • evolve it for the solar age t⊙ • match (Z/X)⊙, L⊙ and R⊙ to better than one part in 105 Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  37. Standard Solar Model Output Information Eight neutrino fluxes: production profiles and integrated values. Only 8B flux directly measured (SNO) so far • Chemical profiles X(r), Y(r), Zi(r) • electron and neutron density profiles (needed for matter effects in neutrino studies) Thermodynamic quantities as a function of radius: T, P, density (r), sound speed (c) Surface helium abundance Ysurf (Z/X and 1 = X + Y + Z leave 1 degree of freedom) Depth of the convective envelope, RCZ Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  38. Standard Solar Model: Internal Structure Temperature Density

  39. Neutrinos from the Sun Helioseismology and the New Opacity Problem

  40. Helioseismology: Sun as a Pulsating Star • Discovery of oscillations: Leighton et al. (1962) • Sun oscillates in > 105 eigenmodes • Frequencies of order mHz (5-min oscillations) • Individual modes characterized by • radial n, angular l and longitudinal m numbers Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  41. Helioseismology: p-Modes • Solar oscillations are acoustic waves • (p-modes, pressure is the restoring force) • stochastically excited by convective motions • Outer turning-point located close to temperature inversion layer • Inner turning-point varies, strongly depends on l • (centrifugal barrier) Credit: Jørgen Christensen-Dalsgaard

  42. Examples for Solar Oscillations + + = http://astro.phys.au.dk/helio_outreach/english/

  43. Helioseismology: Observations • Doppler observations of spectral • lines measure velocities of • a few cm/s • Differences in the frequencies • of order mHz • Very long observations needed. • BiSON network (low-l modes) • has data for  5000 days • Relative accuracy in frequencies • 10-5 Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  44. Helioseismology: Comparison with Solar Models • Oscillation frequencies depend on r, P, g, c • Inversion problem: • From measured frequencies and from a reference solar model • determine solar structure • Output of inversion procedure: dc2(r), dr(r), RCZ, YSURF Relative sound-speed difference between helioseismological model and standard solar model

  45. New Solar Opacities (Asplund, Grevesse & Sauval 2005) • Large change in solar composition: • Mostly reduction in C, N, O, Ne • Results presented in many papers by the “Asplund group” • Summarized in Asplund, Grevesse & Sauval (2005) Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  46. Origin of Changes Spectral lines from solar photosphere and corona • Improved modeling • 3D model atmospheres • MHD equations solved • NLTE effects accounted for in most cases • Improved data • Better selection of spectral lines • Previous sets had blended lines • (e.g. oxygen line blended with nickel line) Meteorites • Volatile elements • do not aggregate easily into solid bodies • e.g. C, N, O, Ne, Ar only in solar spectrum • Refractory elements, • e.g. Mg, Si, S, Fe, Ni • both in solar spectrum and meteorites • meteoritic measurements more robust

  47. Consequences of New Element Abundances What is good • Much improved modeling • Different lines of same element give • same abundance (e.g. CO and CH lines) • Sun has now similar composition • to solar neighborhood New problems • Agreement between helioseismology • and SSM very much degraded • Was previous agreement a coincidence? Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  48. Standard Solar Model 2005: Old and New Opacity Sound Speed Density Adapted from A. Serenelli’s lectures at Scottish Universities Summer School in Physics 2006

  49. Old and New Neutrino Fluxes Bahcall, Serenelli & Basu (astro-ph/0412440 & astro-ph/0511337)

  50. Neutrinos from the Sun Very Low-Energy Solar Neutrinos

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