Propagation and Composition of Ultra High Energy Cosmic Rays Roberto Aloisio INFN – Laboratori Nazionali del Gran Sasso 6th Rencotres du Vientnam Challenges in Particle Astrophysics Hanoi 6-12 August 2006

The spectrum of CR's GZK 2nd Knee Ankle Knee

The first fuzzy picture of the UHECR sky Small angle clustering gives indications about the source number density At most one source in the angular bin of 3 degrees! The sources should have been seen in particular at 1020 eV Correlation??? if no correlations found Bursting Sources??? or High Magnetic Fields??? Exotic Models??? CAVEAT Recently Blasi, De Marco and Olinto found a 5σ inconsistency between the spectrum and the small scale anisotropies measured by AGASA. Using many realizations (MC) As nS decreases (fixed flux) sources become brighter with an increase in clustering probability ns  10-5 Mpc-3 40 sources within 100 Mpc (~20 degrees between sources) Blasi & De Marco (2003) Kachelriess & Semikoz (2004)

Chemical Composition Similar limits from AGASA No conclusive observations at energies E>1018 eV Hires, HiresMIA, Yakutsk proton composition Fly’s Eye, Haverah Park, Akeno mixed composition Fly's Eye [Dawson et al. 98] Transition from heavy (at 1017.5 eV) to light composition (at ~1019 eV) Haverah Park [Ave et al. 2001] No more than 54% can be Iron above 1019 eV No more than 50% can be photons above 4 1019 eV Similar limits from AGASA Proton composition at E>1018 eV not disfavored by experimental observations Hires elongation rate

UHE Proton energy losses protons CMB Universe size 1000 Mpc Adiabatic losses Universe expansion log10[ latt (Mpc)] 100 Mpc Pair production p   p e+ e- Photopion production p   p 0  n  log10[ E (eV)]

particular effect on the flux Pair production A   A e+ e- UHE Nuclei energy losses Universe size Iron helium Universe size Pair production has no particular effect on the flux Photodisintegration A   (A-1) N  (A-2) 2N Depletion of the flux Iron E  1020 eV Helium E 1019 eV

Protons propagation in Intergalactic Space Blasi, De Marco, Olinto (2003) Continuum Energy Losses Protons lose energy but do not disappear. Fluctuations in the pγ interaction start to be important only at E>5  1019 eV. Uniform distribution of sources the UHECR sources are continuously distributed with a density ns. Discrete sources the UHECR sources are discretely distributed with a spacing d. Injection spectrum number of particles injected at the source per unit time and energy Modification factor γg > 2 injection power law Lp source luminosity ns (d) source density (spacing) model parameters Jpunm(E) only redshift energy losses Jp(E) total energy losses

UHECR proton spectrum DIP GZK The energy losses suffered by protons leave their imprint on the spectrum p+ p + e+ + e– DIP GZK p+ N + p sources distribution (mainly GZK) injection spectrum (mainly DIP) way of propagation (magnetic fileds) These features depend on

Proton Dip Experimental evidence of the Dip Best fit values: γg = 2.7 Akeno AGASA Best fit values: γg = 2.7 <Lp> = 4  1043 erg/s ns = 3  10-5 Mpc-3 Berezinsky et al. (2002-2005)

Robustness and Caveats Protons in the Dip come from large distances, up to 103 Mpc. The Dip does not depend on: inhomogeneity, discreteness of sources maximum energy at the source intergalactic magnetic fields (see later...) The interpretation of the DIP in terms of protons pair-production FAILS if: RA, Berezinsky, Grigorieva (2006) heavy nuclei fraction at E>1018 eV larger than 15% (primordial He has nHe/nH0.08) Berezinsky et al. (2004) Allard et al. (2005) RA et al. (2006) the injection spectrum has gg< 2.4

Maximum energy distribution Diffusive shock acceleration tipically shows Maximum energy distribution The maximum acceleration energy is fixed by the geometry of the source and its magnetic field If the sources are distributed over Emax: (β ≈ 1.5) the overall UHECR generation rate has a steepening at some energy Ec (minimal Emax O(1018 eV)) Kachelriess and Semikoz (2005) RA, Berezinsky, Blasi, Grigorieva, Gazizov. (2006)

Energy calibration by the Dip Different experiments show different systematic in energy determination Calibrating the energy through the Dip gives an energy shift E→ λE (fixed by χ2) λAGASA = 0.90 λHiRes = 1.21 λAuger = 1.26 NOTE: λ<1 for on-ground detectors and λ>1 for fluorescence light detectors (Auger energy calibration by the FD)

Intergalactic Magnetic Fields Very poor experimental evidences Faraday rotation Synchrotron and ICS emission Magnetic field concentrated around sources, i.e. in Large Scale Structures No appreciable field in most part of the universe volume Large Scale Structures are characterized by magnetic field produced from compression and twisting of the primordial field Voids are characterized by an appreciable magnetic field Astrophysical sources Cosmological origin Effect of IMF on UHECR deflection diffusion isotropization

Numerical simulations Numerical determination of the IMF is based on LSS and MHD simulations Puzzling results by different groups Sigl, Miniati & Ensslin use an unconstrained simulation putting the observer * close to a cluster Dolag, Grasso, Springel & Tkachev use constrained simulations, being able to reproduce the local Universe High B (100 nG in filaments and 1 nG in voids) High deflection angles: up to 20° at 1020 eV UHECR astronomy nearly impossible Low B (0.1 nG in filaments and 0.01 nG in voids) Low deflection angles: < 1° at 4 1019 eV UHECR astronomy is allowed

Steepening in the flux at The DIP survives also with IMF The IMF effect on the UHE proton spectrum Magnetic Horizon – Low Energy Steepening The diffusive flux presents an exponential suppression at low energy and a steepening at larger energies. Steepening in the flux at E1018 eV 2nd Knee no IMF The DIP survives also with IMF gg=2.7 The low energy cut-off is due to a suppression in the maximal contributing distance (magnetic horizon), its position depends on the IMF. The steepening is independent of the IMF, it depends only on the proton energy losses and coincides with the observed 2nd Knee. The low energy behavior (E<1018 eV) depends on the diffusive regime. RA & Berezinsky (2005) Lemoine (2005)

Galactic Cosmic Rays proton Ep Rigidity models can be rigidity-confinement models or rigidity-acceleration models. The energy of spectrum bending (knee) for nucleus Z is Ez = Z Ep, where Ep » 3×1015 eV is the proton knee. For Iron EFe» 8´1016 eV. Ep proton

EHe helium ?

BUT Kascade data Kascade data 2003: seem to confirm the rigidity model. BUT Kascade data 2005: different results with different Monte Carlo approaches in data reconstruction. Rigidity scenario not confirmed.

Galactic and ExtraGalactic I The Galactic CR spectrum ends in the energy range 1017 eV, 1018 eV. 2nd Knee appears naturally in the extragalactic proton spectrum as the steepening energy corresponding to the transition from adiabatic energy losses to pair production energy losses. This energy is universal for all propagation modes (rectilinear or diffusive): E2K»1018 eV. with IMF without IMF gg=2.7 RA & Berezinsky (2005)

Galactic and ExtraGalactic II Traditionally (since 70s) the transition Galactic-ExtraGalactic CR was placed at the ankle (» 1019 eV). In this context ExtraGalactic protons start to dominate the spectrum only at the ankle energy with a more conservative injection spectrum gg » 2.0  2.2. Problems for the Galactic component Galactic acceleration: Maximum acceleration energy required is very high Emax» 1019 eV Composition: How the gap between Iron knee EFe»1017eV and the ankle (1019 eV) is filled

Conclusions 1. Is there the GZK feature? Auger will soon clarify this point. First results seem to favor the GZK picture. 2. Is there a dip? Spectrum in the range 1018 - 1019 eV could be a signature of proton interaction with CMB (as the GZK feature). 3. Where is the transition Galactic-ExtraGalactic CRs? Precise determination of the mass composition in the energy range 1018 - 1019 eV. Galactic CR (nuclei) at E ≥ 1018 eV challenge for the acceleration of CR in the Galaxy (Emax  EFe) ExtraGalactic CR (protons) at E ≥ 1018 eV discovery of proton interaction with CMB confirmation of conservative models for Galactic CR models for the acceleration of UHECR with γg > 2.4