Stars as Laboratories for Fundamental Physics - MPP Theory Group

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402 Chapter 11 associated with the dissociation of the remaining iron shell it had to work through before it could fly. Recall that the shock wave forms relatively deep inside of the collapsing iron core of the progenitor star. A possible solution is the “delayed mechanism” (Wilson 1983; Bethe and Wilson 1985) where the shock wave lingers at a constant radius for a few 100 ms and then takes off again (Fig. 11.1), powered both by the accretion of material and by the energy deposition of the neutrino flow. In this regard the dilute hot region (ρ of order 10 6 −10 8 g cm −3 , T of order 1 MeV) below the stalled shock plays a major role at absorbing neutrino energy. 63 Unfortunately, it is not certain that the neutrino flux can deposit enough energy to revitalize the shock as the energy transfer is relatively inefficient. However, if one adjusts the amount of neutrino energy transfer to a value above a rather well defined threshold one obtains beautiful explosions (Fig. 11.4). It should be noted that the spectacular explosion of a SN—which at the peak of its lightcurve outshines an entire galaxy—is only a “dirt effect” relative to the release of neutrino energy which equals the gravitational binding energy of the newborn neutron star of about G N M/R ≈ 3×10 53 erg with M ≈ 1.4 M ⊙ and R ≈ 10 km. The total energy released in the kinetic energy of the ejecta and in electromagnetic radiation is a few 64 10 51 erg, on the order of 1% of the total neutron-star binding energy. Therefore, on energetic grounds alone there is no problem at tapping the neutrino flux for explosion energy. A variety of schemes are currently being discussed to achieve a successful shock revival. It is thought that convection may play a major role at transporting energy to the surface of the bloated protoneutron star which has formed after the break-out of the shock (e.g. Burrows and Lattimer 1988; Wilson and Mayle 1988; Mayle and Wilson 1993; Burrows and Fryxell 1992, 1993; Janka and Müller 1993b). It may be that this mechanism can boost the effective neutrino luminosity for a few hundred milliseconds, enough in some calculations to trigger an explosion (see however Bruenn and Mezzacappa 1994). At any rate, in the absence of a fundamental treatment of convection this method is essentially one way of parametrizing the initial amount of neutrino heating below the shock. 63 Goodman, Dar, and Nussinov (1987) proposed that the pair-annihilation process νν → e + e − might be the dominant mode of energy transfer. However, Cooperstein, van den Horn, and Baron (1987) critizised the neutrino emission parameters of that study while Janka (1991) found that a proper treatment of the phase space renders this process less efficient than had been originally thought. 64 This unit is sometimes referred to as 1 foe, for “ten to the fifty one ergs.”
Supernova Neutrinos 403 Convection is also important at transporting energy within the dilute region between the protoneutron star and the stalling shock wave (Herant, Benz, and Colgate 1992; Herant et al. 1994; Janka and Müller 1993a, 1994, 1995a,b; Sato, Shimizu, and Yamada 1993; Burrows, Fig. 11.4. Unsuccessful (upper panel) and successful (lower panel) SN explosion. In each case, the location of several mass shells is shown as a function of time. The thick shaded line indicates the location of the shock. The only difference between the two cases is the adjusted neutrino luminosity from a central source; it was chosen as L ν = 2.10×10 52 erg s −1 (upper panel) and L ν = 2.20×10 52 erg s −1 (lower panel), respectively. In the upper case, which is just below threshold for a successful explosion, the shock displays an interesting oscillatory behavior. (Curves courtesy of H.-T. Janka, taken from Janka and Müller 1993a.)
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Georg G. Raffelt Stars as Laborator
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Table of Contents Preface (xiv) Ack
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Table of Contents vii 4.5 Neutrino
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Table of Contents ix 8. Neutrino Os
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Table of Contents xi 12. Radiative
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Table of Contents xiii 16.3 New Int
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Preface xv Second, particles from d
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Preface xvii search for proton deca
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Preface xix cuts of higher-order gr
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Preface xxi quoted “astrophysical
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Chapter 1 The Energy-Loss Argument
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The Energy-Loss Argument 3 example
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The Energy-Loss Argument 5 trinos,
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The Energy-Loss Argument 7 As a sim
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The Energy-Loss Argument 9 The most
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The Energy-Loss Argument 11 against
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The Energy-Loss Argument 13 ing M;
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The Energy-Loss Argument 15 M ′ (
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The Energy-Loss Argument 17 Table 1
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The Energy-Loss Argument 19 ing to
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The Energy-Loss Argument 21 scalars
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Chapter 2 Anomalous Stellar Energy
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Anomalous Stellar Energy Losses Bou
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Chapter 3 Particles Interacting wit
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Particles Interacting with Electron
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Chapter 4 Processes in a Nuclear Me
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Processes in a Nuclear Medium 119 t
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Processes in a Nuclear Medium 121 f
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Processes in a Nuclear Medium 123 F
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Processes in a Nuclear Medium 129 v
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Processes in a Nuclear Medium 131 4
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Processes in a Nuclear Medium 133 i
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Processes in a Nuclear Medium 135 H
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Processes in a Nuclear Medium 137 I
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Processes in a Nuclear Medium 143 i
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Processes in a Nuclear Medium 145 s
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Processes in a Nuclear Medium 147 E
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Processes in a Nuclear Medium 149 a
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Processes in a Nuclear Medium 151 T
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Processes in a Nuclear Medium 153 r
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Processes in a Nuclear Medium 155 4
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Processes in a Nuclear Medium 157 I
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Processes in a Nuclear Medium 159 t
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Processes in a Nuclear Medium 161 A
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Processes in a Nuclear Medium 163 R
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Chapter 5 Two-Photon Coupling of Lo
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Two-Photon Coupling of Low-Mass Bos
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Chapter 6 Particle Dispersion and D
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Particle Dispersion and Decays in M
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Chapter 7 Nonstandard Neutrinos The
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Nonstandard Neutrinos 253 (“spont
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Nonstandard Neutrinos 255 kinematic
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Nonstandard Neutrinos 257 95% CL ma
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Nonstandard Neutrinos 259 Fig. 7.2.
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Nonstandard Neutrinos 261 In three-
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Nonstandard Neutrinos 263 Flavor mi
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Nonstandard Neutrinos 267 The quant
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Nonstandard Neutrinos 269 dipole mo
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Nonstandard Neutrinos 271 component
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Nonstandard Neutrinos 275 moments.
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Nonstandard Neutrinos 277 7.5.2 Spi
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Nonstandard Neutrinos 279 where m
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Neutrino Oscillations 281 and hence
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Neutrino Oscillations 283 As usual
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Neutrino Oscillations 285 two-flavo
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Neutrino Oscillations 287 Fig. 8.2.
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Neutrino Oscillations 289 8.2.4 Exp
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Neutrino Oscillations 293 Apparentl
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Neutrino Oscillations 297 when the
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Neutrino Oscillations 299 If the ne
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Neutrino Oscillations 301 8.3.5 The
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Neutrino Oscillations 303 It is als
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Neutrino Oscillations 305 system, h
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Neutrino Oscillations 307 say, betw
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Neutrino Oscillations 309 1991). Th
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Oscillations of Trapped Neutrinos 3
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Chapter 10 Solar Neutrinos The curr
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Solar Neutrinos 343 Fig. 10.1. Sola
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Solar Neutrinos 345 There exist vas
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Solar Neutrinos 347 10.2 Calculated
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Solar Neutrinos 349 Fig. 10.3. Norm
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Radiative Particle Decays 453 corre
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Radiative Particle Decays 455 is no
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Radiative Particle Decays 457 Fig.
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Radiative Particle Decays 461 12.3.
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Radiative Particle Decays 463 emiss
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Radiative Particle Decays 465 range
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Radiative Particle Decays 469 weak
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Radiative Particle Decays 471 value
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Radiative Particle Decays 473 ignor
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Radiative Particle Decays 479 Table
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Radiative Particle Decays 481 In or
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Radiative Particle Decays 485 while
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Radiative Particle Decays 487 h 100
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Radiative Particle Decays 491 still
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Chapter 13 What Have We Learned fro
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What Have We Learned from SN 1987A
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Axions 525 Table 14.1. Particle mag
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Axions 527 or axion decay constant.
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Axions 529 The Yukawa coupling h is
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Axions 531 14.2.3 Pseudoscalar vs.
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Axions 533 Notably, the Higgs field
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Axions 535 otic heavy quarks: only
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Axions 537 Various axion models dif
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Axions 539 Bounds on the Yukawa cou
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Axions 541 Fig. 14.5. Astrophysical
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Axions 543 decay into axions. This
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Chapter 15 Miscellaneous Exotica St
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Miscellaneous Exotica 547 Table 15.
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Miscellaneous Exotica 549 15.2.3 Pr
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Miscellaneous Exotica 551 Most rece
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Miscellaneous Exotica 553 Fig. 15.1
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Miscellaneous Exotica 555 A violati
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Miscellaneous Exotica 557 nuclei it
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Miscellaneous Exotica 559 mass, and
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Miscellaneous Exotica 561 backgroun
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Miscellaneous Exotica 563 SN 1987A
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Miscellaneous Exotica 565 which led
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Miscellaneous Exotica 567 For a nar
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Neutrinos: The Bottom Line 569 the
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Neutrinos: The Bottom Line 571 Neut
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Neutrinos: The Bottom Line 573 Apar
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Neutrinos: The Bottom Line 575 is d
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Neutrinos: The Bottom Line 577 of t
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Neutrinos: The Bottom Line 579 siti
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Units and Dimensions 581 In the ast
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Appendix B Neutrino Coupling Consta
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Appendix C Numerical Neutrino Energ
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Numerical Neutrino Energy-Loss Rate
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Numerical Neutrino Energy-Loss Rate
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Appendix D Characteristics of Stell
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Characteristics of Stellar Plasmas
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References 607 Akhmedov, E. Kh., an
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References 609 Bahcall, J. N., Kras
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References 611 Bilenky, S. M., and
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References 613 Cannon, R. D. 1970,
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References 615 Cooper-Sarkar, A. M.
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References 617 Dodelson, S., Friema
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References 619 Fukuda, Y., et al. 1
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References 621 Goyal, A., Dutta, S.
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References 623 Hoogeveen, F., and S
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References 625 Kawakami, H., et al.
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References 627 Landau, L. D., and P
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References 629 Mayle, R. W., and Wi
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References 631 Nieves, J. F., and P
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References 633 Pochoda, P., and Sch
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References 635 Ross, J. E., and All
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References 637 Shlyakhter, A. I. 19
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References 639 Totsuka, Y. 1993, Ne
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References 641 Wiezorek, C., et al.
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Acronyms 643 L l.h. l.h.s. ly LMC L
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Symbols 645 dp differential in phas
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Symbols 647 X j Peccei-Quinn charge
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Subject Index A α-Ori 189 A665, A1
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Subject Index 651 Coulomb gauge 204
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Subject Index 653 H Hayashi line 32
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Subject Index 655 multiple scatteri
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Subject Index 657 neutrino radiativ
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Subject Index 659 Primakoff process
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Subject Index 661 SN 1987A bounds (
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Subject Index 663 supernova core bi
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