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Cosmic Neutrino Background (Cvb): Properties and Detection Perspectives

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Cosmic Neutrino Background (Cvb): Properties and Detection Perspectives Cosmic Neutrino Background (CvB): properties and detection perspectives GianpieroGianpiero ManganoMangano INFN,INFN, SezioneSezione didi Napoli,Napoli, ItalyItaly Relic neutrino production and decoupling ν ν ↔ν ν 1 MeV ≤ T ≤ mµ α β α β νανβ ↔νανβ T = T = T - - ν e γ ναe ↔ναe + - νανα ↔e e Neutrino decoupling As the Universe expands, particle densities are diluted and temperatures fall. Weak interactions become ineffective to keep neutrinos in good thermal contact with the e.m. plasma Rough, but quite accurate estimate of the decoupling temperature Rate of weak processes ~ Hubble expansion rate 2 8πρ R 2 5 8πρ R ν Γ w ≈ σ w v n , H = 2 → G F T ≈ 2 → Tdec ≈ 1 MeV 3M p 3M p ± Since νe have both CC and NC interactions with e Tdec(νe) ~ 2 MeV Tdec(νµ,τ) ~ 3 MeV Neutrino and Photon (CMB) temperatures At T~me, electron- positron pairs annihilate e+e- → γγ heating photons but not the decoupled neutrinos 1/3 Tγ ⎛11 ⎞ = ⎜ ⎟ T ⎝ 4 ⎠ 1 ν fν (p,T) = ep/Tν +1 Neutrino and Photon (CMB) temperatures At T~m , Photon temp falls e slower than 1/a(t) electron- positron pairs annihilate e+e- → γγ heating photons but not the decoupled neutrinos 1/3 Tγ ⎛11 ⎞ = ⎜ ⎟ T ⎝ 4 ⎠ 1 ν fν (p,T) = ep/Tν +1 Neutrinos decoupled at T~MeV, keeping a 1 fν (p,T) = spectrum as that of a relativistic species ep/Tν +1 • Number density d 3 p 3 6ζ (3) n = f (p,T ) = n = T 3 ν ∫ ( 2π)3 ν ν 11 γ 11π 2 CMB • Energy density ⎧ 2 4/3 7π ⎛ 4 ⎞ 4 Massless ⎪ ⎜ ⎟ TCMB 3 120 11 d p ⎪ ⎝ ⎠ ρ = p2 + m2 f (p,T ) → ν i ∫ ν i 3 ν ν ⎨ ( 2π) ⎪ m n Massive m >>T ⎪ ν i ν ν ⎩⎪ Neutrinos decoupled at T~MeV, keeping a 1 fν (p,T) = spectrum as that of a relativistic species ep/Tν +1 • Number density 3 d p 3 -36ζ (3) 3 nAtν = presentf ν112(p,T (ν ν ) = + ν ) ncm γ = per TflavourCMB ∫ ( 2π)3 11 11π 2 • Energy density ⎧ 2 4/3 7π 2 ⎛ 4 ⎞ −54 Massless Ω⎪ ν h ⎜=1.7⎟×10TCMB 3 ⎪120 ⎝11⎠ Contribution2 to 2thed energyp ⎪ ρν = p + mν fν(p,Tν ) → ⎨ m densityi ∫ of the Universei ( 2π)3 ∑ i Massive Ω⎪ h2 =m ni ν ν i ν ⎪ 94.1 eV mν>>T ⎩⎪ Relativistic particles in the Universe At T<me, the radiation content of the Universe is 2 2 ⎡ 4/3 ⎤ π 4 7 π 4 7 ⎛ 4 ⎞ ρr = ργ + ρν = Tγ +3× × Tν = ⎢1+ ⎜ ⎟ 3⎥ργ 15 8 15 ⎣⎢ 8 ⎝11⎠ ⎦⎥ Relativistic particles in the Universe At T<me, the radiation content of the Universe is Effective number of relativistic neutrino species Traditional parametrization of the energy density stored in relativistic particles # of flavour neutrinos:Nν = 2.984 ± 0.008 (LEP data) Extra relativistic particles • Extra radiation can be: scalars, pseudoscalars, sterile neutrinos (totally or partially thermalized, bulk), neutrinos in very low-energy reheating scenarios, relativistic decay products of heavy particles… • Particular case: relic neutrino asymmetries Constraints from BBN and from CMB+LSS CvB details + - At T~me, e e pairs annihilate heating photons e+e- → γγ … and neutrinos. Non thermal features in v distribution (small effect). Oscillations slightly modify the result fν=fFD(p,Tν)[1+δf(p)] ⎡ 2 ⎤ M 8 2GF ()i∂t − Hp∂ p ρ = ⎢ − 2 E , ρ⎥ + C ( ρ ) ⎣⎢ p mW ⎦⎥ -2 νe 10 effect νµ,τ δfx102 p2 e p/T + 1 Results γ γ (%) (%) (%) N Tfin /T0 δρνe δρνµ δρντ eff Instantaneous 1.40102 0 0 0 3 decoupling SM 1.3978 0.94 0.43 0.43 3.046 +3ν mixing 1.3978 0.73 0.52 0.52 3.046 (θ13=0) +3ν mixing 2 1.3978 0.70 0.56 0.52 3.046 (sin θ13=0.047) Dolgov, Hansen & Semikoz, NPB 503 (1997) 426 G.M. et al, PLB 534 (2002) 8 G.M. et al, NPB 729 (2005) 221 Effect of neutrinos on BBN 1. Neff fixes the expansion rate during BBN 8ππ ρ H = 80 3 60 D 4 40 ρ(Neff)>ρ0 →↑He 3 7 He 20 Li ⎛ 4/3 ⎞ 7 ⎛ 4 ⎞ eff ρ = ρ + ρ + ρ = ⎜1+ N v ⎟ρ 0 R γ v x ⎜ ⎜ ⎟ ⎟ γ ⎝ 8 ⎝11⎠ ⎠ -20 4He -40 0 2 4 6 8 2. Direct effect of electron neutrinos and antineutrinos on the n-p reactions BBN: allowed ranges for Neff n /n η = B γ ≅ 274Ω h2 10 10−10 B Using 4He + D data (95% CL) +1.4 Neff = 3.1 −1.2 CvB details Fermi-Dirac spectrum with temperature T and chemical potential µν=ξvTv nν ≠ nν Raffelt 3 n − n 1 ⎛T ⎞ = ν ν = ⎜ ν ⎟ π 2 ξ + ξ 3 Lν ⎜ ⎟ []ν ν nγ 12ζ ( 3 ) ⎝Tγ ⎠ ⎡ 2 4 ⎤ 15 ⎛ ξν ⎞ ⎛ ξν ⎞ ∆ρν = ⎢2 ⎜ ⎟ + ⎜ ⎟ ⎥ More radiation 7 ⎣ ⎝ π ⎠ ⎝ π ⎠ ⎦ Combined bounds BBN & CMB-LSS Degeneracy direction (arbitrary ξe) 8 Hansen et al 2001 Hannestad 2003 6 4 eff 2 N − 0.01≤ ξ ≤ 0.22 ξ ≤ 2.4 ∆ + e µ,τ 0 -2 -0.2 0 0.2 0.4 In the presence of flavor oscillations ? ξe Evolution of neutrino asymmetries BBN Effective flavor equilibrium Dolgov et al 2002 ξ ≤ 0.07 Wong 2002 (almost) established → ν Abazajian et al 2002 − 0.05 ≤ ξ ≤ 0.07 Serpico & Raffelt 2005 µν/Tv very small (bad for detection!) BBN, CMB (LSS) + oscillations 0.4 Cuoco et al 04 0.2 PLANCK e µ ξ ν/Tv + 0 -0.2 0.018 0.02 0.022 0.024 0.026 0.028 0.03 ωb µν/Tv Hamann, Lesgourgues and GM 08 Departure from equilibrium? Pastor, Pinto Raffelt 2008 CvB for optimists V produced by decays at some cosmological epoch Φ → νν Early on: Thermal FD spectrum (TBBN> T>TCMB) Distortion from Φ decay Cuoco, Lesgourgues, GM and Pastor ‘05 Late (T<< TCMB): Unstable DM (e.g. Majoron) Γ Γ Ωv < Ωdm < 0.1 H0 H0 Lattanzi and Valle ’07 Lattanzi, Lesgourgues, GM and Valle (in progress) Present constraints from CMB (WMAP+ACBAR+VSA+CBI) and LSS (2dFGRS+SDSS) + SNIa data (Riess et al.) Model: standard ΛCDM + nonthermal v’s Cl and P(k) computed using CAMB code (Lewis and Challinor 2002) Likelihoods (using COSMOMC Lewis and Bridle 2002)) ΛCDM ΛCDM+R ΛCDM+NT 2 χmin 1688.2 1688.0 1688.0 10 ln[10 Rrad]3.2±0.13.2±0.13.2±0.1 ns 0.97±0.02 0.99±0.03 1.00±0.03 ωb 0.0235±0.0010 0.0231±0.0010 0.0233±0.0011 3 3.2 3.4 3.6 0.95 1 1.05 1.1 0.02 0.0220.0240.026 0.15 0.2 0.25 10 ω ω ωdm 0.121±0.005 0.17±0.03 0.17±0.03 ln [10 R ] n rad b dm θ 1.043±0.005 1.033±0.006 1.033±0.006 τ 0.13±0.05 0.13±0.06 0.15±0.07 β 0.46±0.04 0.48±0.04 0.48±0.04 m0 (eV) 0.3±0.20.8±0.50.7±0.4 Neff 3.04 6±26±2 q 111.25±0.13 h 0.67±0.02 0.76±0.06 0.76±0.05 Age (Gyr) 13.8±0.212.1±0.912.1±0.8 1.02 1.03 1.04 1.05 0.1 0.2 0.3 0.4 0.5 0.6 0.5 1 1.5 2 . ± . ± . ± . θ τ β m ΩΛ 0 68 0 03 0 67 0 03 0 67 0 03 0 zre 14±416±518±6 σ8 0.76±0.06 0.77±0.07 0.77±0.07 TABLE I: Minimum value of the effective χ2 (defined as −2lnL,whereL is the likelihood function) and 1σ confi- dence limits for the parameters of the three models under consideration. The first ten lines correspond to our basis of 4 6 8 10 1 1.2 1.4 1.6 1.8 0.05 0.1 0.15 10 20 30 independent parameters, while the last five refer to related N q = ω (93.2 eV / m ) y parameters. eff ν 0 A * CvB for pessimists A neutrinoless Universe? Beacom, Bell and Dodelson 2004 Models where v’s interact with light (pseudo)scalar particles L = hijviv jφ + gijviv jφ + h.c. vv ↔ φφ couplings < 10-5 from several data (meson decay, 0ββ, SN) For the tightly coupled regime v density strongly reduced, v’s play no role in LSS Delay in matter domination epoch, different content in relativistic species after v decays (Neff=6.6 after decays) Beacom, Bell and Dodelson 2004 Bell, Pierpaoli and Sigurdson 2005 Hannestad 2005 Including CMB in the analysis: •No free streaming (no anisotropic stress) leads to smaller effects on LSS (for massive v’s) change of sub-horizon perturbations at CMB epoch •Change of sound speed and equation of state of the the titghly coupled v - φ fluid •v decays larger Neff i.e. larger ISW effect for CMB smaller effects on LSS (no v left at LSS formation epoch Massless interacting Massive decaying Massive decaying Interacting v-DM Usual picture of Dark Matter: cold collisionless massive particles which decoupled around the weak scale for freeze-out of annihilation processes Ex: neutralino in MSSM with mass of O(100 GeV) Difficulties: Excess of small scale structures Far more satellite galaxies in the Milky Way than observed (from numerical simulation) DM in the MeV range: SPI spectrometer on the INTEGRAL satellite observed a bright 511 KeV gamma line from the galactic bulge Boehm, Fayet and Silk 2003 ψψ → e+e− Framework: light (MeV) DM interacting with neutrinos Several options for lagrangian density G.M., Melchiorri, Serra, Cooray and Kamionkowski 2006 Effects on cosmological scales: if DM - v’s scatterings at work during LSS formation we expect an oscillating behavior in the power spectrum (analogous to baryon – photon fluid during CMB epoch) Equation for velocity perturbations Effect depends upon the parameter Q mF ∫ mDM mF =mDM Bounds on differential opacity Q from SDSS data mDM ∫ mF scenario almost ruled out.
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