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Angelica de Oliveira-Costa

The Cosmic Microwave Background:. New Challenges. Angelica de Oliveira-Costa. University of Pennsylvania. XI Advanced School of Astrophysics. Campos do Jordao, September 2002. The Importance of CMB Polarization:.

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Angelica de Oliveira-Costa

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  1. The Cosmic Microwave Background: New Challenges. Angelica de Oliveira-Costa University of Pennsylvania XI Advanced School of Astrophysics Campos do Jordao, September 2002

  2. The Importance of CMB Polarization: 1. Polarization measurements can substantially improve accuracy with which parameters are measured by breaking the degeneracy between certain parameter combinations. 2. It also offers an independent test of the basic assumptions that underly the standard cosmological model.

  3. Where does CMB Polarization comes from(Hu & White 1997)? CMB polarization is induced via Thomson scattering, either at decoupling or during a later epoch of reonization. The level of polarization induced is linked to the local quadrupole anisotropy of radiation incident on the scattering eletrons. The level of polarization is expected to be 1%-10% of the amplitude of the temperature anisotropies. Important things to know(Kamionkowski et al. 1997, Zaldarriaga 1998): Under coordinate transformations, the Q and U maps transform into a “vector” field on the celestial sphere described by the quantities E and B. E and B can correlate with each other, and with the temperature T. By parity, <EB> and <TB> are zero, <TE> has the largest signal, <EE> is smaller, and <BB> should be zero (except for the cases of gravity-waves present in the last scattering or the existence of polarized foregrounds). TT TE TB TE EE EB TBEB BB

  4. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  5. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  6. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM (wd = Wd h2) 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  7. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos (fn =WHDM/WT) 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  8. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  9. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  10. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  11. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  12. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  13. Our Cosmological Model: Polarization Movies: Matter Buget 1. g = Photons 2. wb = Baryons (H,He,…) 3. wd = CDM 4. fu = Neutrinos 5. WL = Lambda (Dark Energy) 6. Wk = Curvature Input Fluctuations 7. As = Scalar Normalization 8. At = Tensor Normalization 9. ns = Scalar Tilt 10. nt = Tensor Tilt Gastrophysics 11. t = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. www.hep.upenn.edu/~angelica/polarization.html

  14. Princeton IQU Experiment (PIQUE): Ground based experiment (roof of Jadwin Hall). FWHM = 0.235o (100<l<600). Operates at 90 (and 40) GHz. Scans a ring of radius 1o around the NCP (144 pixels). HEMT correlation receiver. Sensitivity ~ 3mK Expected foregrounds < 0.5mK. Team: D. Barkats J. Gundersen M. Hedman S. Staggs B. Winstein Part of analysis effort: A. de Oliveira-Costa M. Tegmark M. Zaldarriaga Hedman et al. (2001)

  15. PIQUE Analysis: We compute 5 power spectra T,E,B,TE & TB with a QE method, and later complement it with the Likelihood analysis. 50 bands w/ dl=20 till l=1000 Netterfield et al. (1997): dTSK =50mK Headman et al. (2001): dTE <14mK 211 (+294,-146) dTB <13mK 212 (+229,-135) dTE(dTB=0) <10mK de Oliveira-Costa et al (2002): dTTE <17mK dTTB <20mK dTEB ??? To do better we need reduce PIQUE pixel noise.

  16. Polarization Observations of the Large Angular Regions (POLAR): Ground based experiment (Madison, WI). FWHM=7o (2<l<20). Operates at 30 GHz Scans at fixed d=43o (300 pixels). HEMT correlation receiver. Expected sensitivity ~ 1-5mK. O’Dell (2002) Team: B. Keating C. O’Dell A. Polnarev J. Steinberger P. Timbie Part of analysis effort: A. de Oliveira-Costa M. Tegmark Keating et al. (2001)

  17. POLAR Results: 3 bands w/ dl=10 till l=30 Smoot et al. (1992): dTDMR =20mK Keating et al. (2001): (Normalized Likelihood Contours) dTE <10mK dTB <10mK dTE(dTB=0) < 8mK de Oliveira-Costa et al (2002): (Band power estimates - same results when average the bands) dTTE <13mK dTTB <11mK dTEB < 4mK O’Dell, Ph.D Thesis (2002)

  18. “Leakage”: Tegmark &de Oliveira-Costa et al. (2001). B2002, l=20: B2002, l=70: 1. E and B are symmetric: There are equal leakage from E to B and vise-versa. 2. Leakage drops with l (E/B separation works well for l>>dl). 3. Map-shape is important: The narrowest dimension of Circle, l=70: B2002, l=20 (disentangle): the map is the limiting factor. 4. Sensitivity is negligible compared with sky coverage: In a situation where sample variance is dominant, this tends to make windowns slightly lobsided. 5. There is no leakage between T & TE and E & TE. 6. There is no leakage between TE & TB, E & EB and B & EB: de Oliveira-Costa et al. (2002). 7. Leakage between E & B can be completed removed: Bunn et al. (2002).

  19. Balloon Observations Of Millimetric Extragalactic Radiation ANd Geophysics (BOOMERanG): Ballon experiment (two 10 day flight). FWHM=10’ (50<l<1000). de Bernardis et al. (2000) Operates between 150 to 450 GHz. 1st flight: 80 & 800(o)2. 2nd flight: 80 & 800(o)2. Bolometers. Sensitivity ~ 7mK (“small regions”) and ~22mK otherwise. Team: UCSB: J. Ruhl, K. Coble, T. Montroy, E. Torbet Caltech: A. Lange, B. Crill, V. Hristov, B. Jones, K. Ganga, P. Manson JPL: J. Bock U.Mass: P. Mauskopf U.Penn: A. de Oliveira-Costa, M. Tegmark U.Toronto: B. Netterfield U.La Sapienza: P. de Bernardis, S. Masi, F. Piacentini, F. Scaramuzzi, N. Vittorio IROE: A. Boscareli Queen Mary: P. Ade

  20. Microwave Anisotropy Probe (MAP): Frequencies(GHz): 22 30 40 60 90 FWHM(o): 0.93 0.68 0.53 0.35 0.23 Sensitivity: ~35mK (all channels & 0.3ox 0.3o pixels) Detector: Differential Radiometer (with polarization) Data release: Jan 2003! Data from 1st full sky scan More info at: http://map.gsfc.nasa.gov

  21. Other CMB Polarization Experiments: Experiment FWHM n(GHz) Receiver Sensitivity Area Site CapMap 3’ 30,90 HEMT 0.2mK 3(o)2 Princeton (300<l<2000) CBI 3’-6o 30 Interferometer 3 of 100(o)2 Atacama (2<l<2000) DASI 10-15’ 30 HEMT 10(o)2 SP (100<l<900) Polatron 2.5’ 100 Bolometer 11mK 5313(’)2 OVRO (300<l<2000) • RoPE 2o 9 HEMT 5mK 560(o)2 LBNL (2<l<50) Compass 15’ 30,40&90 HEMT 8mK U.Wisc. (l<650) BOOMERanG 10’ 150,250&450 Bolometer 7mK,22mK 80-800(o)2 SP (50<l<1000) Maxipol10’ 140&420 Bolometer1.4mK NM MAP 13-41’ 30,40,60&90 HEMT 19mK All sky Space L2 (l<600) Planck-LFI 14’&10’ 70&100 HEMT 6mK All sky Space L2 (l<1500) Planck-HFI 8’&5’ 143&217 Bolometer 6mK All sky Space L2 SPOrt 7o 22,32,60&90 HEMT 80% sky Space Station (2<l<20)

  22. Small Scale CMB Experiments: We propose a Center for High Resolution CMB studies (CfHRC). This center will develop a Millimeter Bolometer Camera (MBC) which will be implemented in the Atacama Cosmology Teslescopy (ACT). Ground based experiment at Atacama desert, Chile. Operates at frequencies 145, 225 & 265 GHz. FWHM=1.7, 1.1 & 0.93’. Scans only in azimuth with the ability to cross-link elevations. Sensitivity/pixel ~ 2, 8 & 16mK (64 nights of quality data). Team: Princeton: N. Jarosik, R. Lupton, L. Page, U. Seljak, D. Spergel, S. Staggs, D. Wilkinson U.Penn: A. de Oliveira-Costa, M. Devlin, B. Jain, M. Tegmark Haverford: S. Boughn, B. Partridge Rutgers: A. Kosowsky U.Toronto: B. Netterfield NASA/GSFC: H. Moseley NIST: K. Irwin

  23. CfHRC Goals: Measure the primary anisotropy beyond the MAP & Planck resolution limits. Find galaxy clusters at z<1 through SZ effect. Measure the amplitude of the CMB gravitational lensing, and therefore probe the mass power spectrum at 1Mpc scales at z~1-2. Detect signature of reonization at z~10 through Vishniac effect. Find all extragalactic mm-wave point sources in 200(o)2 to a sensitivity of 1mJy.

  24. Galactic Microwave Emission: Objectives 1. Accurate modeling and subtraction of Galactic foreground contamination in order to correct measure the CMB power spectrum. 2.Unique opportunity to understand the Galactic emission processes between 10 to 103 GHz. Smoot et al. (1992)

  25. Galactic Microwave Emission: Objectives 1. Accurate modeling and subtraction of Galactic foreground contamination in order to correct measure the CMB power spectrum. 2.Unique opportunity to understand the Galactic emission processes between 10 to 103 GHz. Smoot et al. (1992)

  26. Synchrotron Emission:

  27. Dust Emission:

  28. Free-Free Emission: Reynolds et al. (2001)

  29. COBE/DMR Smoot et al. (1992) At 31GHz we expect DIRBE traces free-free.

  30. Saskatoon OVRO result (Leitch et al. 1997) is much higher than expected for a free-free component.

  31. 19 GHz Spinning dust grains predicts a turn-over at lower frequencies (Draine & Lazarian 1998).

  32. QMAP

  33. Tenerife Jones (1999) Smoking gun: evidence for a turn-over and WHAM correlations only at b<15o.

  34. Galactic Microwave Emission: Objectives 1.Accurate modeling and subtraction of Galactic foreground contamination in order to correct measure the CMB power spectrum. 2.Unique opportunity to understand the Galactic emission processes between 10 to 103 GHz. Smoot et al. (1992)

  35. Galactic Microwave Emission: Objectives 1.Accurate modeling and subtraction of Galactic foreground contamination in order to correct measure the CMB power spectrum. 2.Unique opportunity to understand the Galactic emission processes between 10 to 103 GHz. Smoot et al. (1992)

  36. QMAP Foregrounds:

  37. QMAP Power Spectrum:

  38. Polarized Foregrounds: Residual foregrounds after cleaning 5 MAP channels:

  39. Conclusions: Our ability to measure cosmological parameters using the CMB will only be as good as our understanding of the microwave foregrounds. CMB Polarization is likely to be a goldmine of cosmological information, allowing improved measurements of many cosmological parameters and numerous important cross-checks and tests of the underlying theory. CMB Small Angular Scale maps enables new fundamental cosmological tests.

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