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

Cosmology Overview: The Hot Big Bang Model: 1. Expansion. 2. Large-scale homogeneity & isotropy. 3. Primordial nucleosynthesis. 4. CMB.

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.

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. 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, and are zero, has the largest signal, is smaller, and should be zero (except for the cases of gravity-waves present in the last scattering or the existence of polarized foregrounds). Important things to know (Kamionkowski et al. 1997, Zaldarriaga 1998) : TT TE TB TE EE EB TB EB BB

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM (  d =  d h 2 ) 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos (f n =  HDM /  T ) 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

Matter Buget 1. g = Photons 2.  b = Baryons (H,He,…) 3.  d = CDM 4. f u = Neutrinos 5.  L = Lambda (Dark Energy) 6.  k = Curvature Input Fluctuations 7. A s = Scalar Normalization 8. A t = Tensor Normalization 9. n s = Scalar Tilt 10. n t = Tensor Tilt Gastrophysics 11.  = Reonization Optical Depth From Combinations of Parameters 12. h = Hubble Constant Others Foregrounds, Topology, Defects, etc. Our Cosmological Model: Polarization Movies:

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

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

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

POLAR Results: O’Dell, Ph.D Thesis (2002) Keating et al. (2001):  T E <10  K  T B <10  K  T E (  T B =0) < 8  K de Oliveira-Costa et al (2002):  T TE <13  K  T TB <11  K  T DMR =20  K Smoot et al. (1992):  T EB < 4  K 3 bands w/ dl=10 till l=30 (Normalized Likelihood Contours) (Band power estimates - same results when average the bands)

“Leakage”: Tegmark & de Oliveira-Costa et al. (2001). 1. E and B are symmetric: 2. Leakage drops with l (E/B 3. Map-shape is important: 4. Sensitivity is negligible variance is dominant, this There are equal leakage from E to B and vise-versa. separation works well for l>>dl). The narrowest dimension of the map is the limiting factor. compared with sky coverage : In a situation where sample tends to make windowns slightly lobsided. B2002, l=20:B2002, l=70: Circle, l=70:B2002, l=20 (disentangle): 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).

Balloon Observations Of Millimetric Extragalactic Radiation ANd Geophysics (BOOMERanG): Ballon experiment (two 10 day flight). Operates between 150 to 450 GHz. FWHM=10’ (50<l<1000). Bolometers. Sensitivity ~ 7  K (“small regions”) and ~22  K otherwise. 2 nd flight: 80 & 800( o ) 2. de Bernardis et al. (2000) 1 st flight: 80 & 800( o ) 2. 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

Foregrounds are: Syn, Free-Free, dust, rot.dust & PtS Boomerang Performace:

Microwave Anisotropy Probe (MAP): Frequencies(GHz): FWHM(  ): Sensitivity: ~35  K (all channels & 0.3  x 0.3  pixels) Detector: Differential Radiometer (with polarization) More info at: Data release: Jan 2003! Data from 1st full sky scan

Other CMB Polarization Experiments: Experiment FWHM n(GHz) Receiver Sensitivity Area Site Polatron 2.5’ 100 Bolometer 11  K 5313(’) 2 OVRO †RoPE 2 o 9 HEMT 5  K 560( o ) 2 LBNL Compass 15’ 30,40&90 HEMT 8  K U.Wisc. MAP 13-41’ 30,40,60&90 HEMT 19  K All sky Space L2 Planck-LFI 14’&10’ 70&100 HEMT 6  K All sky Space L2 Planck-HFI 8’&5’ 143&217 Bolometer 6  K All sky Space L2 SPOrt 7 o 22,32,60&90 HEMT 80% sky Space Station (300<l<2000) (2<l<50) (l<650) (l<600) (l<1500) (2<l<20) BOOMERanG 10’ 150,250&450 Bolometer 7  K,22  K ( o ) 2 SP (50<l<1000) Maxipol 10’ 140&420 Bolometer 1.4  K NM DASI 10-15’ 30 HEMT 10( o ) 2 SP (100<l<900) CBI 3’-6 o 30 Interferometer 3 of 100( o ) 2 Atacama (2<l<2000) CapMap 3’ 30,90 HEMT 0.2  K 3( o ) 2 Princeton (300<l<2000)

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). Operates at frequencies 145, 225 & 265 GHz. Ground based experiment at Atacama desert, Chile. Sensitivity/pixel ~ 2, 8 & 16  K (64 nights of quality data). FWHM=1.7, 1.1 & 0.93’. Scans only in azimuth with the ability to cross-link elevations. Team: Haverford: S. Boughn, B. Partridge U.Penn: A. de Oliveira-Costa, M. Devlin, B. Jain, M. Tegmark Princeton: N. Jarosik, R. Lupton, L. Page, U. Seljak, D. Spergel, S. Staggs, D. Wilkinson Rutgers: A. Kosowsky U.Toronto: B. Netterfield NASA/GSFC: H. Moseley NIST: K. Irwin

CfHRC Goals: Measure the primary anisotropy beyond the MAP & Planck resolution limits. Measure the amplitude of the CMB gravitational lensing, and therefore probe the mass power spectrum at 1Mpc scales at z~1-2. Find galaxy clusters at z<1 through SZ effect. 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.

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 10 3 GHz. Smoot et al. (1992)

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 10 3 GHz. Smoot et al. (1992)

Quantifying Galactic Emission in a CMB data:

Synchrotron Emission:

Dust Emission:

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

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

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

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

QMAP

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

Frequency Dependence for 4 Latitude Slices: Colors are for DIRBE, Haslam & WHAM

IRAS images from Cloud MBM20: A simple visual comparison of these images suggests that although the large scale features match up, small scale features can be quite different. Therefore spinning dust should be traced by shorter wavelenght dust maps.

Dust Correlations for the mm DIRBE Maps:

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 10 3 GHz. Smoot et al. (1992)

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 10 3 GHz. Smoot et al. (1992)

QMAP data analysis: We introduced new methods for removal of 1/f-noise and scan- synchronous offsets.

Boomerang & Maxima. Other experiments: Xu et al. (2001)

Tegmark & Efstathiou (1996)

QMAP Foregrounds:

QMAP Power Spectrum:

Polarized Foregrounds: Residual foregrounds after cleaning 5 MAP channels:

Conclusions: 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. Our ability to measure cosmological parameters using the CMB will only be as good as our understanding of the microwave foregrounds.