Supernova Neutrinos: Challenges and Prospects

Transcription

1 Supernova Neutrinos: Challenges and Prospects Sanjay Reddy, Institute for Nuclear Theory, Seattle Introduction to supernova - a nuclear physics perspective Supernova neutrinos - why they are so relevant Neutrino interactions - how neutrino s see dense matter The neutrino sphere - correlated matter at low density Charged current reactions Neutron star tomography Outlook.

2 Inevitability of Collapse of Massive Stars Successive stages of nuclear burning produces a inert core. The inner core is supported by electrondegeneracy pressure. It becomes unstable to gravitational collapse at the Chandrashekar mass.

3 Core Collapse Supernova 1500 km Core collapse t collapse ~100 ms 3X10 7 km 10 km Shock wave E shock ~10 51 ergs Quasi-static ~ 1 s 100 km neutrinos diffuse out of the dense newly born neutron star to carry away almost all of the ~ 3 x ergs

4 ergs = MeV Neutrinos Past:! SN 1987a: ~ 20 neutrinos in support of supernova theory Supernova Neutrinos Future:! Can detect ~10,000 neutrinos from galactic supernova Pons et al. (2002)

5 A Neutron Star is Born hot mantle: T~5-40 MeV n~ n0 Ye~ shock wave core: T~10-20 MeV n~1-2 n0 Ye~0.3 R=10 km neutrino sphere R=100 km R=150 km

6 Protoneutron Star Evolution Gravitational binding energy is stored as internal energy: Thermal and neutrino degeneracy energy.! Typical timescale for energy and lepton number loss in the core is s. 40s 55s Pons et al. (1999) t=1s 10s 20s

7 Protoneutron Star Evolution Neutrino diffusion drives the evolution. Star contracts and heats up during deleptonization. Interior cooling begins at t ~ 15 s. Cooling of the outer regions occurs at early times. Pons et al. (1999) Neutrino diffusion and possibly convection sets the time scale.

8 Supernova Neutrino Emission n~ n0! Neutrino interactions here determine the energy spectrum Neutrino interactions here determine the time scale core: n~1-4 n0 R=10 km R=15 km <ε> (MeV) neutrino sphere Figure: adapted from Roberts 2012 ν e ν e ν x Time (s)

9 Neutrinos couple to matter through neutral and charged current 1 = Neutrino Interactions Z d 3 k 3 (2 ) 3 R(E 1,E 3, cos ) f 3 (1 f 1 ) L cc int = G F 2 l µ j µ W for ν l + B 2 l + B 4 L nc int = G F 2 l ν µ jµ Z for ν l + B 2 ν l + B 4 R(E 3,E 1, cos ) f 1 (1 f 3 ) scattering-in scattering-out pair-production +R(E 1, E 3, cos ) (1 f 1 )(1 f 3 ) R( E 1,E 3, cos ) f 1 f 3 pair-annihilation q 0 = E 1 + E 3 Dense Matter q 0 = E 1 E 3 q 0 = E 1 E 3! pair scattering q

10 Correlations & Neutrino Scattering Neutrinos see more than one particle in the medium. Nature of spatial and temporal correlations between nuclei, nucleons and electrons affect the scattering rate. E3,k3 E1,k1 q0 = E1- E3 q= k1- k3 R At small q0 and q the neutrino cannot resolve single particles. Sawyer (1975, 1989) Iwamoto & Pethick (1982) Horowitz & Wherberger (1991) Raffelt & Seckel (1995)

11 Neutrinos Scatter From Fluctuations of Density & Spin E1,k1 E3,k3 q0 = E1- E3 q= k1- k3 Differential Scattering/Absorption Rate: d (E 1 ) = G2 F d cos dq (E 1 q 0 ) 2 (1 + cos ) SV RPA (q 0,q)+(3 cos ) SA RPA (q 0,q) spectrum of density fluctuations spectrum of spin fluctuations

12 The Camera Analogy E3,k3 E1,k1 The neutrino takes pictures of the plasma q0 = E1- E3 q= k1- k3 shutter speed 1/aperture

13 n > p=e Ground State & Fluctuations Low temperature and density: uniform matter one large neutron-rich! nucleus + electrons several small neutron-rich! nuclei + electrons nuclear attraction Coulomb repulsion High temperature and density: n > p=e Entropy or Pauli exclusion and Nuclear repulsion

14 Nuclear pasta At intermediate temperature and density complex structured phases are likely: ρ= g/cm3 and T=1-3 MeV Figures from: Schneider, Horowitz, Hughto & Berry (2013)

15 Composition and fluctuations in the neutrino sphere: ρ= g/cm 3 and T=3-8 MeV How does the competition between density, entropy and nuclear binding energy resolve under these conditions? This is the reason we are here at the ECT workshop. ρ= g/cm 3 T=3-8 MeV F=E-TS What is the free energy difference between these configurations

16 Kc(α), f (1984).to of the Guggenheim systematics, in a manner analogous the calculation, mass fractions cannot be acprevious treatments [75, 79]. 22. M. Schmidt, G.11 Ro pke, and10 H.7 Schulz, Ann. Data Phys. ( 10 Symmetric matter termined. To avoid this problem we have proexperimental Tests 202, 57 (1990). L-S180 equilibrium constants foratcluster formation T(ρ)257, from133 Expt. 4.3 Test of the nuclear matter EOS low densities 23. J. P. Bondorf, etbe al., Phys. Rept. (1995). 10 nstead of mass fractions. In contrast to mass 10 L-S220 Can we reliably predict and 24. G. Ro pke, Phys. Rev. C 79, (2009). With our confidence in the temperature and density deter6 luster formation equilibrium constants, such Rev. as C 71, S.inTypel, (2005). minations bolstered the by the consistency exhibited Fig. 5, Phys. L-S375 measure reduced we have addressed various aspects of clustering in the low 26. G. Ro pke, N.-U. Kla hn, S. T particle formation, i.e., In theoretical 109 Bastian, D. Blaschke, T.Typel density nuclear matter produced. models et al. NS effective nuclear binding? and H. H. Wolter, Nucl. Phys. A897, 70 (2013). cluster mass fractions are commonly used to characterize the degree of clusterization in low Howpke etb32 al. Q 27.matter. D. Vautherin and D. M. Brink, Phys.RoLett. density 5 8 ever, di erent theoretical models include various di erent c (1970); Phys.10Rev. C 5, 626 (1972). 2to di erences 2 competing species. This leads in particuh. Shen et al. te larand symmetry energy of low density nuclear matter 5 n p mass fractions quoted. If all relevant28. species not G.are Ro pke, A. Grigo, K. Sumiyoshi, and Hong Shen, G. Shen et al. included in the calculation, mass fractions cannot be ac7 Part. Nucl. Lett. 2, Data (2005). curately determined. To avoid et thisal.problem Typel, Ropke (2010)we have pron p L-S180 G. Shen et al.a 29.formation G. Ro pke, M. Schmidt and H. 4Schulz, Nucl. Phys. posed that equilibrium constants for cluster be employed instead of mass fractions. In contrast to (1984). mass L-S Schwenk-Hor 10 fractions, cluster formation equilibrium constants, such as L-S A. Kolomiets et al., Phys. Rev. C 55, 1376 (1997). that for particle formation, i.e., Typel et al. NSE n K ( ) = n n (39) Shift of Binding Energies of Lig Kc(α), fm9 10 n, and n are the particle, neutron and equilibsities, respectively, should be independent of 10 onfined tion and choice of competing species. Voskresenska here thereference [39] depicts 6 from a comparison be31. S. Shlomo and V. M. Kolomietz, Rep. Prog. Phys. Ropke et al. QSM n 5 (39) c ( ) = Hempel and S cted. (2005). H. Shen et al. RMF-TM1 experimentallykderived constants nn np equilibrium (2011). G. Shen et al. RMF-NL3 32. G. Ro pke, Nucl. of Phys. A867, rformed resulting from models employing a variety where n, nn, and np are the particle, neutron and G. Shen et al. RMF-FSUGold G. Ro pke proton respectively, should be33. independent of et al., 104 Phys. Stat. Sol. (b) 100, 215 (1980 gystate as adensities, of proposed for astrophysical applications Schwenk-Horowitz Viral 0 EOS proton fraction and choice of competing species. 34. V. Baran, et al., Phys. Rept. 410, 335 (2005).0.01 Voskresenskaya et al. g(rmf) ). Typ16]. Figure 6 from reference [39] depicts a comparison be35. M. B. G.R., TsangPRC et3 al.,79, Phys. Rev. (2009) Lett. 102, (2 Hempel and Schaffner-Bielich NSE (Excl. Vol)+RMF tween our experimentally derived equilibrium constants 10 Hartreeprisingly the calculated values ofwuenschel the equilib36. a S. etetal., Phys. Rev. C 79, (R) (20 and those resulting from models employing variety of S. Typel al., equations for astrophysical tendof state to proposed converge at the lower densities. rant Dirac37. applications C. Sfienti, et al.,0 Phys. Rev. 102, Lett (2009 [8, 9, 10, 11, 16]. PRC 81, (2010) 3 Comparison, nuc/fm ehelowest densities sampled, however, J. equilibwangthere et al., are Phys. Fig. Rev. C6.ρ75, (2007).of ex uninot surprisingly the calculated values38. of the those from various cal rium constant tend tohigher converge atdensities, the 39. lowerl. densities. Qin et al., Phys. Rev. Lett. 108, EOS (2012). di erences. At 0.01 to 0.03 approxg.r., NP A 867, 66 (2011) Even at the lowest densities sampled, however, there are Fig. 6. Comparison of experimental values of K ( ) with c 40. R. Wada et al., Phys. Rev. C 85, (2012).

17 The energy difference between neutrons and protons in neutron-rich matter is related to symmetry energy S( ) Correlations play a key role in increasing at densities and temperatures relevant to the neutrino sphere. Nuclear Symmetry Energy S( ) Nuclear & Neutron Matter E( n, p) =E( )+S( ) symmetry energy per baryon U sym [MeV] n p 2 + S( )=a V a + L 3 (u 1) u = o T =10, NSE T = 5, NSE T = 2, NSE T =10, QS T = 5, QS T = 2, QS T =10, RMF T = 5, RMF T = 2, RMF baryon density n [fm -3 ] Hagel, Natowitz, Ropke et al. (2014)

18 Is the binding relevant for supernova neutrinos?! For charged current reactions difference between neutron and proton binding is relevant. Final state electron blocking is strong for the reaction: e e q 0 U = U n U p Large q0 crucial to overcome blocking n p Reddy, Prakash & Lattimer (1998) Roberts (2012) Martinez-Pinedo et al. (2012) Roberts & Reddy (2012) In neutron rich matter protons are deeply bound. Extent of binding depends on correlations and clusters.

19 Charged Current Reactions at Low Density dγ cos θde e = G2 F 2π p e E e (1 f e (E e )) [ (1 + cos θ)s τ (q 0,q)+g 2 A (3 cos θ)s στ(q 0,q) ] Energy shift helps overcome electron final state blocking. Enhances νe absorption Larger energy needed to produce neutrons suppresses anti-νe absorption. 1/λ (m -1 ) λ νe /λ νe e-05 1e-06 1e ν e + n -> e - + p ν e + p -> e + + n Bremsstrahlung T = 8 MeV n B = 0.02 fm -3 Y e = (µ eq =0) ε ν (MeV) Roberts, Reddy & Shen (2012)

20 B Neutrino Spectra are Sensitive to Symmetry Energy < > (MeV) e (U = 0) e x e (U=RMF) 6 Roberts, Reddy & Shen (2012) Time (s) Time evolution of electron neutrino spectrum could be a useful diagnostic

21 Implications of modifying = Large favors production of neutron-rich nuclei in supernova.! Large drives collective neutrino flavor oscillations closer to the neutron star.! Large should be detectable and its time evolution can constrain the low density symmetry energy.

22 Convection and equation of state Convection is driven by composition gradients.! Buoyancy of matter depends on is composition through the density dependence of the symmetry energy.! Low L drives convection and large values suppresses it.! Roberts, Cirigliano, Pons, Reddy, Shen, Woosley (2012)! convection neutrino diffusion

23 Observable signatures of convective transport Neutrino flux is enhanced during convection.! There is break in the light curve (when convection ends).! Fraction of events between 3-10 s provides good discrimination. Count Rate (s 1 ) w/o convection Counts (3 s > 10 s)/ Counts (0.1 s > ) Counts (0.1 s > 1 s)/ Counts (0.1 s > ) Time (s) No Convection g =0.6 GM3 Convection MF GM3 Convection g =0.6 GM3 Convection g =0.6 IU-FSU Large L (stiff) Small L (soft) Count rate in Super-Kamiokande for galactic supernova at 10 kpc. Roberts, Cirigliano, Pons, Reddy, Shen, Woosley (2012)

24 Neutron Star Tomography core: n~4 n0 T ~ 5-10 MeV R~12 km neutrino sphere At late stages t > 20 s the neutrino sphere moves inward exposing higher density matter. The fluctuations in the pasta will become relevant. Neutrino diffusion at the highest densities now sets the timescale. Final rapid plunge of neutrino luminosity could reveal the most exciting information about neutron star interiors.

25 Hypothetical weakly interacting particles Since neutrinos are trapped for ~20 s other light particles with weaker coupling to matter can change the cooling timescale. E other > ergs/g/s -shortens the neutrino signal by a factor of 2 Axions R 2 =0.44 µm R 2 =1.4 µm KK Gravitons Raffelt (1996) Pons et al (2001)

26 Outlook Neutrino interactions at supra and sub nuclear densities shape supernova neutrino emission. Neutrino cross-sections at sub-nuclear density are likely to be affected by correlations and clustering. Charged current rates are key to understanding spectral differences between different neutrino types. Sensitive to nuclear symmetry energy at low density. Models of warm low density matter remain poorly constrained by experiment. Detailed modeling supernova neutrino emission can unravel deep mysteries. Good chance your grandchildren will be proud.