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THE DARK SIDE
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  1. THEDARKSIDE Openproblemswith Dark Matter & Dark Energy A review and tentative solutions Silvio Bonometto Dep. Physics G.Occhialini Milano-Bicocca LNGS – September 2005

  2. In collaboration with Roberto Mainini Loris Colombo Dip.Fisica G.Occhialini Milano-Bicocca Andrea Maccio’ Theor.Physics University of Zurich Alessandro Gardini Illinois Giuseppe Murante Torino Sebastiano Ghigna ? ? Elena Pierpaoli Caltech Stefano Borgani Trieste ……….

  3. Problems with DM • WHY there must be (non-barionic) DM • Halo numerical profiles (NFW) • also vs observed profiles • Galaxy satellite abundance • Problems with DE • WHY there must be DE • Fine tuning problems (56 o.m.) • Coincidence problems • In particular: • Why DE density ~ DM density just today ?

  4. A twofold aim • Describing current problems, • namely those related to LSS • (2) Introducing “dual axion” model • Same # of parameters as SCDM SCDM : these param- eters Cosmological Parameters in all models N(photons)/N(baryons) Density/Critical Density CDM Density/Baryon density baryogenesis geometry rm ~ rb LCDM cosmology 1 extra parameter: Matter density/Critical density Dynamical DE (+ 1 parameter) Coupled DE (+2 parameter)

  5. Underlying ideology: Should astrophysics put limits on extra parameters new physics discovered and constrained Microph. solution Even better: (new) physics requires DE & DM, setting their parameters in the fair range: A MICROPHYSICAL WAY OUT FROM FINE TUNING & COINCIDENCE • An alternative view (Kolb,Riotto,Matarrese,…2005 • see also Buckert 1980, Ellis 1990 …) • gmn=hmn+hmn • standrd hmn defined by a(t) & k coming • from assigned state eqs. (p=wr) • hmn initially linear, then developing • non-linearities • new hmn initially linear • when extreme non-linear. developed • backg.state eqn modified Geometr. solution

  6. Standard approach to cosmology covariant form conformal time smoothing

  7. in space, not in time ! Stress-energy tensor pseudoconservation

  8. Stress-energy pseudoconservation eqn can be solved without knowing a(t)

  9. Scalar field lagrangian

  10. for DE comp. DE, at early times, just subdominant Dominantr : radiat. or DM fulfilling equation r = wp Major problem with Dark Energy Consistent with CMB data Brax & Martin, 1999, 2001 Brax, Martin & Riazuelo 2000

  11. A further possibility concerning scalar field Exploit its phase… New potential term V1(q,f) New eqs of motion Potential energy also for q COUPLING NO PRESSURE DM

  12. A short discussion of these classical argument

  13. Milky Way reconstruction OPTICAL RADIUS SPIRAL ARMS CORE BRILLIANCE EQUATES THE NIGHT SKY notice the bar Spiral arms: density waves in the EXPONENTIAL DISK Optical radius ro ~ 3l ROTATION CURVE v(r) increasing almost up to optical radius Keplerian region expected 21cm radio data : NO FALL up to ~7-8 times the optical radius Dark Matter mass up to 7-8 times (or more) the visible mass

  14. From ngal and mass in Stars, dust, gas, etc: Wb ~ 0.01 a=2 M prop R If 21cm rotation curve extends ~10 times opt rad then M ~ 10 Moptical IF TOTAL MASS IS 10 TIMES BARYONIC MASS BBNS FORBIDS ITS BEING BARYONIC

  15. e.g. you don’t give M or R for small galaxies, but their m.s.v. Virial theorem extensively used both to analyse cosmic objects and to compare then with theoretical or numerical predictions USE IT NOW FOR GALAXY CLUSTERS

  16. Large galaxy cluster, optical image M ~ 10^15 Msun Notice strong lensing -> Mcl In agreement with mass from virial equilibrium ~10 times baryon mass in gal. + X-ray emitting hot gas Wm > 0.10-0.15

  17. Spherical top-hat fluctuat. growth vs. a(t): analyt. solution in SCDM cosm. Cluster form by action of pure gravity Energy conservation + virial equilibrium Rfinal=0.5 Rtop (SCDM) Top expansion : density contrast ~15 Virialization : density contrast ~180 Further contraction requires heat dissipation. This happens when galaxies form. Dissipative forces have no role in cluster shaping

  18. Understanding CMB and deep sample data Not unique value, but d distribution (normal)

  19. Adiabatic & isocurvature (isothermal) More details in next pages spectra

  20. Understanding CMB and deep sample data Linear Fluctuations : At early times or over very large scale today Coupling C(j) (also C=0)

  21. baryons : Holding when photons and baryons are strictly coupled, or when coupling just begins to fade photons : massless neutrino Fluctuations: either gravitionallu unstable, or free streaming Expansion in harmonics is an essential ingredient of any numerical calculation of anisotropy and polarization of CMB; it is unessential to compute transfer functions (see below)

  22. RADIATION & BARYON SOUND WAVES WHILE DM FLUCTUATIONS INCREASE (AFTER EQUIV.) Different phases in waves, when recombination occurs, will be visible in tranf.funct. and Cl DM fluctuations grow all the time, but between their entry in horizon and equivalence, if they occur in this order

  23. first available whole sky data angular resolution 7 degrees

  24. first release of WMAP data notice increased resolution, there is however a good correspondence between COBE & WMAP

  25. A short review on the approach followed in order to compare CMBR data with models Stockes parameters

  26. Expansion in scalar spherical harmonics Expansion in spin 2 spherical harmonics 4p rotation periodicity

  27. Temperature fluctuations T-E cor- relation E-mode polarization LCDM How spectra depend on optical depth

  28. C~4b/mp SUGRA potential vs WMAP

  29. Best fit param. for LCDM cosm. IMPORTANT TO OUTLINE THAT THESE VALUES ARE TRUE WHEN ASSUMING A LCDM COSMOLOGY

  30. all percents rather similar (not so great) SUGRA as good as LCDM

  31. Large t values were main WMAP discovery Range of values depends on the model Also Wtot ~ 1 Wb much less than Wm:Non-baryonic DM needed Wm much less than Wtot : DE needed

  32. Results obtained with MCMC technique 1 and 2-sigma confidence levels notice how l=log(L/GeV) almost unconstrained

  33. A similar plot for 1/f coupled cosmologies notice how better constrained all parameters and L in particular

  34. Varying total matter density Varying baryonic density

  35. Bending due to Meszaros effect Wiggles due to baryon contribution to final amplitude NON-LINEAR EVOLUTION NUMERICAL SIMULATIONS the 2df sample Numerical simulations

  36. 300 Mpc 3Mpc

  37. HIGHEST RESOLUTION SIMULATION OF DM HALO

  38. Slope of halo profiles. A problem? Concentration, Instead, depends on objects size : c = rc /rv c ~ 5-7 (clusters) c ~10 (galaxies) Navarro 2003 central profile slope indipendent on object size Diemand et al 2004/2005 NFW profile (1997) central slope -1 Moore et al 2001 central slope -1.5

  39. Swaters et al 2003 De Block et al 2001 LSB galaxies measuring of the slope of the core profile

  40. How galaxy rot data are fitted co work out central slope Swaters et al 2003 reproducing slope estimate on simulated objects with ~NFW profile Spekkens et al 2005

  41. ANDROMEDA SATELLITES

  42. Moore et al 1999 Klypin et al 1999 Where are the missing galaxy satellites ? 2 solution: missing satellites did not form missing satellites are there, but invisible…

  43. Bullock, Weinberg & Kravtsov 2002 PopIII stars reionize the Universe at z~8 Gas infall in low-mass halos is suppressed after reionization Working only for z(reion) ~ 8 If reionization earlier mechanism fails (Maccio’ et al 2005) t~16 requires z(reion)~18 MECHANISM NEEDED TO REMOVE BARYONS FROM SMALL HALOS

  44. Towards the conclusion … OTHER MECHANISM TO HAVE DM-ENRICHED SATELLITES REQUIRED OTHER EVIDENCES OF DM-BARYON SEGREGATION ON OTHER SCALES e.g. THE L-T CLUSTER PROBLEM IN THIS CASE ASTROPHYSICAL SOLUTIONS PROPOSED BUT THEIR EFFICIENCY IS STILL DISPUTED DM-DE COUPLING PROVIDES MECHANISM FOR BARYON-DM SEGREGATION AN EVIDENCE IN FAVOR OF COUPLING WITHIN THE DARK SECTOR ?

  45. Still many problems in the dark …….