The Particle Universe (continued) Joakim Edsjö Stockholm University Joakim Edsjö Stockholm University

Outline Cosmic rays from our own galaxy Ways to search for dark matter High-energy cosmic rays Cosmic rays from our own galaxy Ways to search for dark matter High-energy cosmic rays

Candidates The main dark matter candidates are: –Baryonic dark matter can only be a small part –Axions could be part of the DM –Neutrinos probably only part of the DM –Weakly Interacting Massive Particles, WIMPs could be a major part of the DM Will focus on these!

WIMP Dark Matter Produced thermally in the early Universe A particle with a weak interaction cross section has   ~ 1. In supersymmetric extensions of the standard model, such particles arise naturally.

MSSM – Mass spectrum

The MSSM – parameters  -Higgsino mass parameter M 2 -Gaugino mass parameter m A -mass of CP-odd Higgs boson tan  -ratio of Higgs vacuum expectation values m 0 -scalar mass parameter A b -trilinear coupling, bottom sector A t -trilinear coupling, top sector

The MSSM – general The Lightest Supersymmetric Particle (LSP) Usually the neutralino. If R-parity is conserved, it is stable. The Neutralino –  Gaugino fraction 1.Select MSSM parameters 2.Calculate masses, etc 3.Check accelerator constraints 4.Calculate relic density < h2 h2 < 0.5 ? 6.Calculate fluxes, rates,... Calculation done with

Relic density – simple approach Decoupling occurs when  < H We have Figure from Jungman, Kamionkowski and Griest, Phys. Rep. 267 (1996) 195.

Relic density – accurate approach Solve the Boltzmann equation –properly taking the thermal average –including the full annihilation cross section (all annihilation channels, thresholds, resonances). –including so called coannihilations between other SUSY particles present at freeze-out.

Relic density vs mass and composition The neutralino is cosmologically interesting for a wide range of masses and compositions!

LEP  h 2 <  h 2 > 1 Low sampling The m  -Z g parameter space Higgsinos Mixed Gauginos

WIMP search strategies Direct detection Indirect detection: –neutrinos from the Earth/Sun –antiprotons from the galactic halo –positrons from the galactic halo –gamma rays from the galactic halo –gamma rays from external galaxies/halos – synchrotron radiation from the galactic center / galaxy clusters Direct detection Indirect detection: –neutrinos from the Earth/Sun –antiprotons from the galactic halo –positrons from the galactic halo –gamma rays from the galactic halo –gamma rays from external galaxies/halos – synchrotron radiation from the galactic center / galaxy clusters

Direct detection - general principles WIMP + nucleus  WIMP + nucleus Measure the nuclear recoil energy Suppress backgrounds enough to be sensitive to a signal, or... Search for an annual modulation due to the Earth’s motion around the Sun

Direct detection – scattering diagrams Spin-independent scattering Spin-dependent scattering + diagrams with gluons Diagrams from Jungman, Kamionkowski and Griest, Phys. Rep. 267 (1996) 195.

Direct detection – example spectra Differential rateGamma background in Ge-detector Figures from Jungman, Kamionkowski and Griest, Phys. Rep. 267 (1996) 195.

Direct detection – current limits Spin-independent scatteringSpin-dependent scattering Direct detection experiments have started exploring the MSSM parameter space!

Direct detection - DAMA claim Annual modulation seen – interpreted as WIMPs However, not seen by other experiments...

Neutralino capture and annihilation Sun  Earth Detector  velocity distribution  scatt  capture  annihilation interactions hadronization int.  int.

Neutrino telescopes – how do they work? The neutrino interacts with a nucleus in the ice and creates a muon. The muon emits Cherenkov radiation. The radiation is recorded by photomultipliers and the muon track can be reconstructed. The neutrino interacts with a nucleus in the ice and creates a muon. The muon emits Cherenkov radiation. The radiation is recorded by photomultipliers and the muon track can be reconstructed.

Neutrino telescopes Capture and annihilation Capture AnnihilationEvaporation (negligible) Evolution equation SolutionDependencies C–f(v),  ,  scatt, composition of Earth/Sun C A –  ann,  (r) in Earth/Sun

Neutrinos and muons from the Earth’s atmosphere  Use the Earth as a filter by looking for upgoing muons.  Only atmospheric neutrinos remain as a background.  Use the Earth as a filter by looking for upgoing muons.  Only atmospheric neutrinos remain as a background.

The Amanda detector

Limits:  flux from the Earth/Sun EarthSun

Dama region and WIMPs from Earth/Sun SunEarth

Annihilation in the halo Neutralinos can annihilate in the halo producing antiprotons positrons gamma rays synchrotron radiation (from e + /e – in magn. fields) neutrinos from the– Earth – Sun Neutralinos can annihilate in the halo producing antiprotons positrons gamma rays synchrotron radiation (from e + /e – in magn. fields) neutrinos from the– Earth – Sun Bergström and Snellman, ’84 Gondolo & Silk, ’00 Silk & Srednicki, ’84; Stecker, Rudaz & Walsh, ’85 Silk & Srednicki, ’84 Freese, ’86; Krauss, Srednicki & Wilczek, ’86 Gaisser, Steigman & Tilav, ’86 Silk, Olive and Srednicki, ’85 Gaisser, Steigman & Tilav, ’86

Gamma rays Monochromatic At one-loop, neutralinos can annihilate to i.e. monochromatic gamma rays. Continuous Neutralinos can also produce a continuum of gamma rays,

Gamma rays Monochromatic At one-loop, neutralinos can annihilate to i.e. monochromatic gamma rays. Continuous Neutralinos can also produce a continuum of gamma rays, Features directionality – no propagation uncertainties low fluxes, but clear signature strong halo profile dependence Features (compared to gamma lines) much lower energy many more gammas per annihilation rather high fluxes, even away from the galactic center not a very clear signature

Gamma signals from the halo The flux of gamma rays in a direction  in a solid angle  is given by MSSM part Halo part

Gamma signal from neutralinos Continuous gammasGamma lines

Gamma lines – rates in GLAST Bergström, Ullio & Buckley, ’97 NFW halo profile, ∆Ω ≈ 1 sr

Gamma lines – rates in ACTs  ZZ Bergström, Ullio & Buckley, ’97 NFW halo profile, ∆Ω ≈ sr

Continuum gammas – fluxes Flux at high galactic latitudes  small halo profile dependence

Gamma fluxes from simulated halo Continuous gammasGamma lines N-body simulations from Calcáneo-Roldan and Moore, Phys. Rev. D62 (2000)

Diffuse extra-galactic gammas Flux of diffuse gammas from neutralino annihilation in external dark matter halos. Moore halo profile. High-rate MSSM models. E -2.1 extrapolation  10  detection E -2.7 extrapolation  20  detection EGRET 86 GeV 166 GeV

Diffusion model of the Milky Way h g ≈0.1 kpc h h ≈3–20 kpc r 0 ≈8.5 kpc R h ≈20 kpc D l ≈D l 0 (1+R/R 0 ) 0.6 D 0 ≈6  cm 2 s -1 R≈p/|Z| R 0 ≈3 GV

Galaxy model Halo profile Energy losses Propagation model The diffusion model with free escape at the boundaries Modified isothermal sphere, Navarro, Frenk and White, Moore et al., etc. Inelastic scattering gives rise to energy losses (included as a ‘tertiary’ source function for antiprotons).

The diffusion equation DiffusionGalactic windEnergy lossSource The diffusion, galaxy and energy loss parameters are derived from cosmic ray studies.

Cosmic ray composition Li, Be and B are overabundant. Sc-Mn are overabundant This overabundance is believed to be due to spallation of C and Fe respectively by interactions with the interstellar medium. Observed abundances  Diffusion equation  Diffusion and galaxy parameters Observations of radioactive isotopes  further constraints on diffusion and galaxy parameters.

Antiproton background Naively, the background below 1 GeV would be very small, but... energy losses p-He interactions reacceleration are all important. Background antiprotons are produced when cosmic rays hit the interstellar medium:

Antiproton signal from neutralinos Antiproton source function Put into the diffusion equation taking the galaxy model into account. The antiprotons meet the solar wind. Take this modulation into account.

Antiproton signal Easy to get high fluxes, but...

Antiprotons – fits to Bess data Background onlyBackground + signal  No need for, but room for a signal.

Positron fluxes from neutralinos Compared to antiprotons, energy losses are much more important essentially only local halo properties are important higher energies due to more prompt annihilation channels (ZZ, W + W -, etc) propagation uncertainties are higher solar modulation uncertainties are higher

Positrons – signal fluxes Compared to antiprotons, the fluxes are typically lower (except at high masses), but...

Positrons – example spectra...the positron spectra can have features that could be detected! The signal strength needs to be boosted, e.g. by clumps, though......and the fit is not perfect

WIMP conclusions There is mounting evidence for dark matter in the Universe. There are many different dark matter candidates. One of the favourite candidates are WIMPs, of which supersymmetric neutralinos have the desired dark matter properties The WIMP rates (direct and indirect) in many different experiments can be high and sometimes have a nice feature to be distinguished from the background. There is mounting evidence for dark matter in the Universe. There are many different dark matter candidates. One of the favourite candidates are WIMPs, of which supersymmetric neutralinos have the desired dark matter properties The WIMP rates (direct and indirect) in many different experiments can be high and sometimes have a nice feature to be distinguished from the background.

About 100 muons/m 2 sec proton muons mesons Cosmic Rays

High-energy cosmic rays Spectrum measured up to ~ (a few)·10 20 eV. Balloon and satellite experiments: below ~10 16 eV Air shower arrays: up to ~10 21 eV eV is about the same energy as a tennis ball from Boris Becker! ~E -2.7 ~E -3.0

Detectors High Resolution Fly’s Eye (HiRes), under construction Auger, under construction Fly’s Eye Agasa + more detectors in tha past (Haverah Park, Yakutsk,...) and on the drawing board (OWL,...)

Pierre Auger Observatory An array of over 1600 particle detectors will measure shower particles as they hit the ground. In addition, during clear, moonless nights, the showers will be viewed as they traverse the atmosphere. Covers an area of about 3000 km 2 (Agasa ~100 km 2 )

Orbiting Wide-angle Light-collectors OWL Monitor km 2 of atmosphere with 10% efficiency. Record about events/year above eV

Extensive Air Showers - Development Number of particles in shower: Multiplication process stops when the energy is  c. Maximum number: Shower maximum at:

Extensive Air Showers Highest energy AGASA event observed: 2·10 20 eV.

Extensive Air Showers - Direction From timing information at the surface, the direction can be obtained.

Do they keep their direction? The gyroradius for a charged particle (charge Ze) in a magnetic field is given by For small deflection angles, For high energies, E~p,

Greisen Zatsepin Kuzmin (GZK) cut-off Consider a cosmic ray proton. At high energy, it can interact with a CMB photon: CM energy is enough to produce final state when Mean free path for eV proton:

Energy spectrum at the ankle AGASA events above eV Where do they come from?

Source location and GZK cut-off Either, the sources are nearby, the production energies are extremeley high, the estimated energies are wrong, or the highest-energies cosmic rays are not protons.

Cosmic Ray acceleration Scientific American, Credit: George Kelvin Cosmic rays are believed to be accelerated in shocks, e.g. around a supernova, black hole etc. Supernova remnants can accelerate up to ~10 16 eV.

Possible sources Highest energy cosmic rays can only be produced outside of our own galaxy