Stars as Laboratories for Fundamental Physics - MPP Theory Group

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276 Chapter 7 The best limits on µ νe are based on the use of reactor neutrinos as a source. The measurement of the ν e -e scattering cross section by Reines, Gurr, and Sobel (1976) was interpreted by Kyuldjiev (1984) to yield a bound of µ νe < 1.5×10 −10 µ B . Since then, the reactor ν e spectrum has been much better understood. Vogel and Engel (1989) stressed that a literal interpretation of the old results by Reines, Gurr, and Sobel (1976) would actually yield evidence for a dipole moment of about 2−4×10 −10 µ B . There is, however, a more recent limit of µ νe < 2.4×10 −10 µ B (7.28) from the Kurchatov Institute (Vidyakin et al. 1992). When interpreting this bound as a limit on the ν e -ν τ transition moment, recall that for an experimentally allowed ν τ mass in the 10 MeV regime this limit would not apply as reactor neutrinos have relatively low energies (below about 10 MeV). Therefore, it remains useful to consider the bounds from the neutrino beam at LAMPF with a ν e endpoint energy of 52.8 MeV µ νe < 10.8×10 −10 µ B , µ νµ < 7.4×10 −10 µ B (7.29) (Krakauer et al. 1990). A similar limit of µ νµ < 8.5×10 −10 µ B was obtained by Ahrens et al. (1990). An improvement of the bound on µ νe by an order of magnitude or more can be expected from a new reactor experiment currently in preparation (“MUNU experiment,” Broggini et al. 1990). However, results will not become available before a few years from now. The transition moments between ν τ and other sequential neutrinos are bounded by the above experiments involving initial-state ν e ’s or ν µ ’s or their antiparticles. The diagonal ν τ magnetic dipole moment, however, is much less constrained because no strong ν τ sources are available in the laboratory; the ν τ -e cross section has never been measured. However, the calculated ν τ flux produced in a proton beam dump from the decay of D s mesons can be used to derive an upper limit on the ν τ -e cross section which can be translated into a bound µ ντ < 5.4×10 −7 µ B , (7.30) (Cooper-Sarkar et al. 1992). Of course, in this range the ν τ electromagnetic cross section would far exceed a typical weak interaction one.
Nonstandard Neutrinos 277 7.5.2 Spin-Flip Scattering in Supernovae Neutrino scattering by photon exchange can also be important in astrophysical settings, notably if neutrinos are Dirac particles. The magnetic or electric dipole coupling is such that it flips the helicity of relativistic neutrinos, i.e. the final state is r.h. for an initial l.h. neutrino. This spin flip is of no importance in experiments where the electron recoil is measured, but it can have dramatic consequences in supernovae where l.h. neutrinos are trapped by the standard weak interactions. The spinflip scattering by a electromagnetic dipole interaction would produce “wrong-helicity” states that could freely escape unless they scattered again electromagnetically. The SN 1987A neutrino signal indicates that this anomalous cooling channel cannot have been overly effective, yielding a constraint of around µ < ν ∼ 3×10 −12 µ B on all Dirac diagonal or transition moments in the sense of Eq. (7.25); see Sect. 13.8.3 for a more detailed discussion. One should keep in mind, however, that for dipole moments in this range the spin precession in the strong macroscopic magnetic fields that are believed to exist in and near SN cores could also cause significant left-right transitions. Notably, the back conversion of r.h. neutrinos could cause a transfer of energy between widely separated regions of the SN core, and might even help at the explosion (Sect. 13.8.3). The role of relatively large neutrino dipole moments in SN physics has not been elaborated in enough depth to arrive at reliable regions of parameters that are ruled out or ruled in by SN physics and the SN 1987A neutrino signal. 7.5.3 Spin-Flip Scattering in the Early Universe Neutrino spin-flip scattering has important consequences in the early universe as it can bring some or all of the “wrong-helicity” Dirac neutrino degrees of freedom into thermal equilibrium. The usual big bang nucleosynthesis (BBN) argument previously mentioned in Sect. 7.1.5 allows one to exclude this possibility because even one additional thermally excited neutrino degree of freedom appears to be forbidden by the spectacular agreement between the predicted and observed primordial light-element abundances. This argument was first advanced by Morgan (1981a,b). Unfortunately, he used an unrealistically small cutoff for the Coulomb divergence of the spin-flip scattering cross section, leading to an overestimate of the efficiency by which r.h. Dirac neutrinos can be brought into ther-
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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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Page 303 and 304: Neutrino Oscillations 283 As usual
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Page 309 and 310: Neutrino Oscillations 289 8.2.4 Exp
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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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Solar Neutrinos 351 a certain range
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Solar Neutrinos 353 While helium an
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Solar Neutrinos 355 Bahcall and Pin
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Solar Neutrinos 357 Xu et al. (1994
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Solar Neutrinos 359 Table 10.4. Spe
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Solar Neutrinos 361 Table 10.5. Pre
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Solar Neutrinos 363 rates attribute
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Solar Neutrinos 365 10.3.4 Water Ch
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Solar Neutrinos 367 Fig. 10.10. Dif
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Solar Neutrinos 369 Fig. 10.13. Mea
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Solar Neutrinos 371 Table 10.9. Cur
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Solar Neutrinos 373 of order the an
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Solar Neutrinos 375 Fig. 10.16. 37
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Solar Neutrinos 377 GALLEX, or Home
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Solar Neutrinos 379 Table 10.10. Me
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Solar Neutrinos 381 deficits in all
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Solar Neutrinos 383 The ν e surviv
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Solar Neutrinos 385 The allowed ran
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Solar Neutrinos 387 10.7 Spin and S
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Solar Neutrinos 389 10.8 Neutrino D
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Solar Neutrinos 391 The small- and
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Solar Neutrinos 393 flux above its
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Chapter 11 Supernova Neutrinos The
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Supernova Neutrinos 397 Fig. 11.1.
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Supernova Neutrinos 399 beyond the
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Supernova Neutrinos 401 ate neutrin
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Supernova Neutrinos 403 Convection
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Supernova Neutrinos 405 Hayes, and
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Supernova Neutrinos 407 11.2 Predic
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Supernova Neutrinos 409 must then b
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Supernova Neutrinos 411 Figure 11.6
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Supernova Neutrinos 413 signal (Sec
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Supernova Neutrinos 415 time) on 23
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Supernova Neutrinos 417 All of the
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Supernova Neutrinos 419 to multiple
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Supernova Neutrinos 421 the final-s
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Supernova Neutrinos 423 ter of 53 e
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Supernova Neutrinos 425 Given these
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Supernova Neutrinos 427 thorough st
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Supernova Neutrinos 429 The forward
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Supernova Neutrinos 431 Another int
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Supernova Neutrinos 433 ber of ν e
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Supernova Neutrinos 435 A “normal
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Supernova Neutrinos 437 Fig. 11.19.
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Supernova Neutrinos 439 The approxi
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Supernova Neutrinos 441 be taken to
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Supernova Neutrinos 443 sense that
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Supernova Neutrinos 445 presence of
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Supernova Neutrinos 447 An even mor
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Chapter 12 Radiative Particle Decay
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Radiative Particle Decays 451 one c
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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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