Probing the Origin of Neutrino Mass and Neutrino Properties
Transcript Of Probing the Origin of Neutrino Mass and Neutrino Properties
Probing the Origin of Neutrino Mass and Neutrino Properties with Supernova Neutrinos †
Ina Sarcevic
University of Arizona
• Models of Neutrino Mass with a Low Scale Symmetry Breaking • New Interactions of Supernova Relic Neutrinos • Probing Neutrino Properties with Supernova Neutrinos • Experimental Detection of the New Interactions via Light Scalar
(HyperK, GADZOOKS, UNO, MEMPHYS) • BBN and SN1987A Constraints on the Parameter Space of the
Neutrino Mass Models
† H. Goldberg, G. Perez and I. Sarcevic, JHEP 11, 023 (2006); J. Baker, H. Goldberg, G. Perez and I. Sarcevic, hep-ph/0630218.
• Observation of neutrino flavor conversion from solar (SuperK, SNO), atmospheric (SuperK) and terrestrial (KamLand, K2K) neutrino data has provided firm evidence for neutrino flavor conversion.
• Recent SuperK data on the L/E-dependence of atmospheric neutrino events and the new spectrum data from terrestrial experiments (KamLand, K2K) is the first time evidence of the expected oscillatory behavior. This strongly favors non-vanishing sub-eV neutrino masses.
Data/Prediction (null osc.)
Ratio
1.8 1.6 1.4 1.2
1 0.8 0.6 0.4 0.2
0 1
2
3
4
10
10
10
10
L/E (km/GeV)
1.4 1.2
1 0.8 0.6 0.4 0.2
0 20
2.6 MeV prompt analysis threshold
KamLAND data best-fit oscillation best-fit decay best-fit decoherence
30
40
50
60
70
80
L0/Eνe (km/MeV)
Neutrino Mass Models with Low Scale of Symmetry Breaking
Chacko, Hall, Okui and Oliver, PR D70 (2004) Chacko, Hall, Oliver and Perelstein, PRL 94 (2005) Davoudiasl, Kitano, Kribs and Murayama, PR D71(2005)
• Consider the effective Lagrangian below EWB scale and close to the neutrino flavor symmetry breaking scale: LDν = Lkin + yν φνN + V (φ) LM ν = Lkin + yν φνν + V (φ)
ν is an active neutrino, V (φ) is the scalar potential (for the global case this contains interaction between φ and the additional Goldstone bosons).
• After the symmetry breaking, neutrino gets the mass =⇒ mν = yν f , where f =< φ > and f << MW .
• Strongest limits on f come from cosmology and astrophysics: light scalars not to be in thermal equilibrium during big bang nucleosynthesis (BBN) gives a limit on f of approximately f ≥ 10 keV.
Chacko et al., PRL 94 (2005) Davoudiasl et al., PRD 71 (2005)
• Similar bound is obtained by demanding that SN cooling not be modified in the presence of these additional fields.
Y. Farzan, PRD 67 (2003)
• Combining BBN bound and assuming that the heaviest neutrino mass, mhν ∼ 4 × 10−2 eV, we find that the strength of the interaction between the scalar and the neutrinos yν is rather weak,
yν ≤ 10−5 ,
New Interactions of Supernova Relic Neutrinos
H. Goldberg, G. Perez and I.S., JHEP 11, 023 (2006).
• Dramatic modification of the supernova neutrino flux (diffuse or burst) through interactions between SN neutrinos and the cosmic background neutrinos:
ν
ν
G
ν
ν
νSN + νCMB → G → νν
• Typical SRN energies are above solar neutrino energies and bellow the atmospheric ones ⇒ likely to be observed by SuperK, GAZOOKS, HyperK, KamLand, UNO and MEMPHYS.
• Unique possibility of detecting the extra light degrees of freedom as well as cosmic background neutrinos!
10 SRN
1
0.1 10
1 0.1
0 20 40 60 80 100
01
The Effect of New Interactions on SRN Flux
• SRN neutrino energies will be redistributed after the interaction
• Significant distortion of the SRN flux as a result of redistribution SRN flux can have regions of depletion relative to flux without new interactions SRN flux can have regions of enhancement relative to flux without new interactions
• These modifications could be detected at large neutrino detectors
SRN Flux with New Interactions
• Depletion of flux in region EνRes/(1 + z) ≤ EνObs ≤ EνRes
• Replenishment of flux from 0 energy back up to Eunscattered for each neutrino energy in resonance region
dF dEΝ cm 2 s 1 MeV 1
1 0.8 0.6 0.4 0.2
5
z 0.5 10 15 20 EΝ MeV
Resonance Including Cosmological Expansion
• Consider the conditions on the coupling for which there is sizable resonant degradation of the original flux of supernova neutrinos.
• Probability that a neutrino, created at red shift z, with energy (1 + z)E arrives unscattered at the detector with energy E is given by:
z
dz¯
P (E, z) = exp − 0 H(z¯)(1 + z¯) n(z¯) σνν→φ(2mν (1 + z¯)E) ,
where n(z¯) = (1 + z¯)3 n0 is the background neutrino density at redshift z¯ and n0 56 cm−3.
• Large depletion of the initial SRN flux if yν >∼ 4 × 10−8 1MkeGV
Accumulative Resonance
• There will be resonant absorption of the original neutrino flux as well as replenishment from neutrinos re-emitted in the decay of a G.
• A neutrino emitted with energy Ei ≥ E∗ from a source at redshift z undergoes resonant scattering at redshift z¯ < z, so that
E∗ = Ei 1 + z¯ 1+z
• Neutrinos from the decay have flat energy distribution, E = f E∗ , 0 ≤ f ≤ 1 , where f varies uniformly over the region [0,1].
• The spectrum with absorption has a dip at E = E∗, and is shifted downward from the spectrum absent resonant absorption. The complete effect of neutrinos emitted with non-resonant energies, passing through resonance, and then replenishing the flux at lower energies, is what we call accumulative resonance.
Ina Sarcevic
University of Arizona
• Models of Neutrino Mass with a Low Scale Symmetry Breaking • New Interactions of Supernova Relic Neutrinos • Probing Neutrino Properties with Supernova Neutrinos • Experimental Detection of the New Interactions via Light Scalar
(HyperK, GADZOOKS, UNO, MEMPHYS) • BBN and SN1987A Constraints on the Parameter Space of the
Neutrino Mass Models
† H. Goldberg, G. Perez and I. Sarcevic, JHEP 11, 023 (2006); J. Baker, H. Goldberg, G. Perez and I. Sarcevic, hep-ph/0630218.
• Observation of neutrino flavor conversion from solar (SuperK, SNO), atmospheric (SuperK) and terrestrial (KamLand, K2K) neutrino data has provided firm evidence for neutrino flavor conversion.
• Recent SuperK data on the L/E-dependence of atmospheric neutrino events and the new spectrum data from terrestrial experiments (KamLand, K2K) is the first time evidence of the expected oscillatory behavior. This strongly favors non-vanishing sub-eV neutrino masses.
Data/Prediction (null osc.)
Ratio
1.8 1.6 1.4 1.2
1 0.8 0.6 0.4 0.2
0 1
2
3
4
10
10
10
10
L/E (km/GeV)
1.4 1.2
1 0.8 0.6 0.4 0.2
0 20
2.6 MeV prompt analysis threshold
KamLAND data best-fit oscillation best-fit decay best-fit decoherence
30
40
50
60
70
80
L0/Eνe (km/MeV)
Neutrino Mass Models with Low Scale of Symmetry Breaking
Chacko, Hall, Okui and Oliver, PR D70 (2004) Chacko, Hall, Oliver and Perelstein, PRL 94 (2005) Davoudiasl, Kitano, Kribs and Murayama, PR D71(2005)
• Consider the effective Lagrangian below EWB scale and close to the neutrino flavor symmetry breaking scale: LDν = Lkin + yν φνN + V (φ) LM ν = Lkin + yν φνν + V (φ)
ν is an active neutrino, V (φ) is the scalar potential (for the global case this contains interaction between φ and the additional Goldstone bosons).
• After the symmetry breaking, neutrino gets the mass =⇒ mν = yν f , where f =< φ > and f << MW .
• Strongest limits on f come from cosmology and astrophysics: light scalars not to be in thermal equilibrium during big bang nucleosynthesis (BBN) gives a limit on f of approximately f ≥ 10 keV.
Chacko et al., PRL 94 (2005) Davoudiasl et al., PRD 71 (2005)
• Similar bound is obtained by demanding that SN cooling not be modified in the presence of these additional fields.
Y. Farzan, PRD 67 (2003)
• Combining BBN bound and assuming that the heaviest neutrino mass, mhν ∼ 4 × 10−2 eV, we find that the strength of the interaction between the scalar and the neutrinos yν is rather weak,
yν ≤ 10−5 ,
New Interactions of Supernova Relic Neutrinos
H. Goldberg, G. Perez and I.S., JHEP 11, 023 (2006).
• Dramatic modification of the supernova neutrino flux (diffuse or burst) through interactions between SN neutrinos and the cosmic background neutrinos:
ν
ν
G
ν
ν
νSN + νCMB → G → νν
• Typical SRN energies are above solar neutrino energies and bellow the atmospheric ones ⇒ likely to be observed by SuperK, GAZOOKS, HyperK, KamLand, UNO and MEMPHYS.
• Unique possibility of detecting the extra light degrees of freedom as well as cosmic background neutrinos!
10 SRN
1
0.1 10
1 0.1
0 20 40 60 80 100
01
The Effect of New Interactions on SRN Flux
• SRN neutrino energies will be redistributed after the interaction
• Significant distortion of the SRN flux as a result of redistribution SRN flux can have regions of depletion relative to flux without new interactions SRN flux can have regions of enhancement relative to flux without new interactions
• These modifications could be detected at large neutrino detectors
SRN Flux with New Interactions
• Depletion of flux in region EνRes/(1 + z) ≤ EνObs ≤ EνRes
• Replenishment of flux from 0 energy back up to Eunscattered for each neutrino energy in resonance region
dF dEΝ cm 2 s 1 MeV 1
1 0.8 0.6 0.4 0.2
5
z 0.5 10 15 20 EΝ MeV
Resonance Including Cosmological Expansion
• Consider the conditions on the coupling for which there is sizable resonant degradation of the original flux of supernova neutrinos.
• Probability that a neutrino, created at red shift z, with energy (1 + z)E arrives unscattered at the detector with energy E is given by:
z
dz¯
P (E, z) = exp − 0 H(z¯)(1 + z¯) n(z¯) σνν→φ(2mν (1 + z¯)E) ,
where n(z¯) = (1 + z¯)3 n0 is the background neutrino density at redshift z¯ and n0 56 cm−3.
• Large depletion of the initial SRN flux if yν >∼ 4 × 10−8 1MkeGV
Accumulative Resonance
• There will be resonant absorption of the original neutrino flux as well as replenishment from neutrinos re-emitted in the decay of a G.
• A neutrino emitted with energy Ei ≥ E∗ from a source at redshift z undergoes resonant scattering at redshift z¯ < z, so that
E∗ = Ei 1 + z¯ 1+z
• Neutrinos from the decay have flat energy distribution, E = f E∗ , 0 ≤ f ≤ 1 , where f varies uniformly over the region [0,1].
• The spectrum with absorption has a dip at E = E∗, and is shifted downward from the spectrum absent resonant absorption. The complete effect of neutrinos emitted with non-resonant energies, passing through resonance, and then replenishing the flux at lower energies, is what we call accumulative resonance.