The Cosmic Background
Imager
Steven T. Myers
National Radio Astronomy Observatory
Socorro, NM
UNM – Oct 14, 2003
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The Cosmic Background Imager
• A collaboration between
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Caltech (A.C.S. Readhead PI)
NRAO
CITA
Universidad de Chile
University of Chicago
• With participants also from
– U.C. Berkeley, U. Alberta, ESO, IAP-Paris, NASA-MSFC,
Universidad de Concepción
• Funded by
– National Science Foundation, the California Institute of
Technology, Maxine and Ronald Linde, Cecil and Sally
Drinkward, Barbara and Stanley Rawn Jr., the Kavli Institute,
and the Canadian Institute for Advanced Research
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The Instrument
• 13 90-cm Cassegrain antennas
– 78 baselines
• 6-meter platform
– Baselines 1m – 5.51m
• 10 1 GHz channels 26-36 GHz
– HEMT amplifiers (NRAO)
– Cryogenic 6K, Tsys 20 K
• Single polarization (R or L)
– Polarizers from U. Chicago
• Analog correlators
– 780 complex correlators
• Field-of-view 44 arcmin
– Image noise 4 mJy/bm 900s
• Resolution 4.5 – 10 arcmin
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3-Axis mount : rotatable platform
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Other Interferometers: DASI, VSA
• DASI @ South Pole
• VSA @ Tenerife
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CBI Instrumentation
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CBI Operations
• Observing in Chile since Nov 1999
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NSF proposal 1994, funding in 1995
Assembled and tested at Caltech in 1998
Shipped to Chile in August 1999
Continued NSF funding in 2002, to end of 2004
• Telescope at high site in Andes
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16000 ft (~5000 m)
Located on Science Preserve, co-located with ALMA
Now also ATSE (Japan) and APEX (Germany), others
Controlled on-site, oxygenated quarters in containers
• Data reduction and archiving at “low” site
– San Pedro de Atacama
– 1 ½ hour driving time to site
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Site – Northern Chilean Andes
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CBI in Chile
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A Theoretical
Digression
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The Cosmic Microwave Background
• Discovered 1965 (Penzias & Wilson)
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2.7 K blackbody
Isotropic
Relic of hot “big bang”
3 mK dipole (Doppler)
• COBE 1992
– Blackbody 2.725 K
– Anisotropies 10-5
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Thermal History of the Universe
Courtesy Wayne Hu – http://background.uchicago.edu
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CMB Anisotropies
• Primary Anisotropies
– Imprinted on surface of “last scattering”
– “recombination” of hydrogen z~1100
– Primordial (power-law?) spectrum of potential fluctuations
• Collapse of dark matter potential wells inside horizon
• Photons coupled to baryons >> acoustic oscillations!
– Electron scattering density & velocity
• Velocity produces quadrupole >> polarization!
– Transfer function maps P(k) >> Cl
• Depends on cosmological parameters >> predictive!
– Gaussian fluctuations + isotropy
• Angular power spectrum contains all information
• Secondary Anisotropies
– Due to processes after recombination
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Acoustic Oscillations
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Power Spectrum of the CMB
Courtesy Wayne Hu – http://background.uchicago.edu
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Dependence on Geometry
Courtesy Wayne Hu – http://background.uchicago.edu
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Dependence on Baryon content
Courtesy Wayne Hu – http://background.uchicago.edu
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Effects of Damping
Courtesy Wayne Hu – http://background.uchicago.edu
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Secondary Anisotropies
Courtesy Wayne Hu – http://background.uchicago.edu
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Gravitational Secondaries
• Due to CMB photons passing through potential
fluctuations (spatial and temporal)
• Includes:
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Early ISW (decay, matter-radiation transition at last scattering)
Late ISW (decay, in open or lambda model)
Rees-Sciama (growth, non-linear structures)
Tensors (gravity waves, ‘nuff said)
Lensing (spatial distortions)
Courtesy Wayne Hu – http://background.uchicago.edu
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Scattering Secondaries
• Due to variations in:
– Density
• Linear = Vishniac effect
• Clusters = thermal Sunyaev-Zeldovich effect
– Velocity (Doppler)
• Clusters = kinetic SZE
– Ionization fraction
• Coherent reionization suppression
• “Patchy” reionization
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2ndary SZE Anisotropies
• Spectral distortion of CMB
• Dominated by massive halos (galaxy
clusters)
• Low-z clusters: ~ 20’-30’
• z=1: ~1’  expected dominant signal in
CMB on small angular scales
• Amplitude highly sensitive to s8
A. Cooray (astro-ph/0203048)
P. Zhang, U. Pen, & B. Wang (astro-ph/0201375)
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Seven Pillars of the CMB
(of inflationary adiabatic fluctuations)
Minimal
Inflationary
parameter
set
•Large Scale Anisotropies
•Acoustic Peaks/Dips
•Damping Tail
•Gaussianity
Quintessence
Tensor fluc.
Broken Scale
Invariance
•Secondary Anisotropies
•Polarization
•Gravity Waves
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Images of the CMB
WMAP Satellite
BOOMERANG
ACBAR
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After WMAP…
• Power spectrum
– measured to l < 1000
– Primary CMB
– First 3 peaks
Courtesy Wayne Hu – http://background.uchicago.edu
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…and Planck
• Power spectrum
– measured to l < 1000
– Primary CMB
– First 6 peaks
Courtesy Wayne Hu – http://background.uchicago.edu
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CMB Interferometry
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Interferometers
• Spatial coherence of radiation pattern contains
information about source structure
– Correlations along wavefronts
• Equivalent to masking parts of a telescope aperture
– Sparse arrays = unfilled aperture
– Resolution at cost of surface brightness sensitivity
• Correlate pairs of antennas
– “visibility” = correlated fraction of total signal
• Fourier transform relationship with sky brightness
– Van Cittert – Zernicke theorem


j
.
2

.
ul

vm
V (u, v)   I (l , m)  e
dl.dm
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The Fourier Relationship
• The aperture (antenna) size smears out the
coherence function response
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Like a double-slit experiment with widening slits
Interference plus diffraction pattern
Lose ability to localize wavefront direction = field-of-view
Small apertures = wide field
• An interferometer “visibility” in the sky and Fourier
planes:
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The uv plane and l space
• The sky can be uniquely described by spherical
harmonics
– CMB power spectra are described by multipole l ( the angular
scale in the spherical harmonic transform)
• For small (sub-radian) scales the spherical harmonics
can be approximated by Fourier modes
– The conjugate variables are (u,v) as in radio interferometry
– The uv radius is given by l / 2
• The projected length of the interferometer baseline
gives the angular scale
– Multipole l = 2 B / l
• An interferometer naturally measures the transform of
the sky intensity in l space
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Interferometry of the CMB
• An interferometer “visibility” in the sky and Fourier
planes:
• The primary beam and aperture are related by:
CMB
peaks
smaller
than this !
CBI:
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Mosaicing in the uv plane
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Power Spectrum and Likelihood
• Statistics of CMB (Gaussian) described by power spectrum:
Break into bandpowers
Construct covariance
matrices and perform
maximum Likelihood
calculation:
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CBI Beam and uv coverage
• 78 baselines and 10 frequency channels = 780
instantaneous visibilities
– Frequency channels give radial spread in uv plane
• Pointing platform rotatable to fill in uv coverage
– Parallactic angle rotation gives azimuthal spread
– Beam nearly circularly symmetric
• Baselines locked to platform in pointing direction
– Baselines always perpendicular to source direction
– Delay lines not needed
– Very low fringe rates (susceptible to cross-talk and ground)
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Calibration and Foreground Removal
• Calibration scale ~5%
– Jupiter from OVRO 1.5m (Mason et al. 1999)
– Agrees with BIMA (Welch) and WMAP
• Ground emission removal
– Strong on short baselines, depends on orientation
– Differencing between lead/trail field pairs (8m in RA=2deg)
– Use scanning for 2002-2003 polarization observations
• Foreground radio sources
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Predominant on long baselines
Located in NVSS at 1.4 GHz, VLA 8.4 GHz
Measured at 30 GHz with OVRO 40m
Projected out in power spectrum analysis
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Power Spectrum Estimation
• Method described in Paper IV (Myers et al. 2003)
• Large datasets
– > 105 visibilities in 6 x 7 field mosaic
– ~ 103 independent
• Gridded “estimators” in uv plane
– fast!
– Not lossless, but information loss insignificant
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Construct covariance matrices for gridded points
Maximum likelihood using BJK method
Output bandpowers
Wiener filtered images constructed from estimators
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The Computational Problem
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Tests with mock data
• The CBI pipeline has been extensively tested using
mock data
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Use real data files for template
Replace visibilties with simulated signal and noise
Run end-to-end through pipeline
Run many trials to build up statistics
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Wiener filtered images
• Covariance matrices can be applied as Wiener filter
to gridded estimators
• Estimators can be Fourier transformed back into
filtered images
• Filters CX can be tailored to pick out specific
components
– e.g. point sources, CMB, SZE
– Just need to know the shape of the power spectrum
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Example – Mock deep field
Raw
Noise
removed
CMB
Sources
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CBI Results
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CBI 2000 Results
• Observations
– 3 Deep Fields (8h, 14h, 20h)
– 3 Mosaics (14h, 20h, 02h)
– Fields on celestial equator (Dec center –2d30’)
• Published in series of 5 papers (ApJ July 2003)
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Mason et al. (deep fields)
Pearson et al. (mosaics)
Myers et al. (power spectrum method)
Sievers et al. (cosmological parameters)
Bond et al. (high-l anomaly and SZ) pending
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CBI Deep Fields 2000
Deep Field Observations:
•3 fields totaling 4 deg^2
•Fields at d~0 a=8h, 14h, 20h
•~115 nights of observing
•Data redundancy  strong
tests for systematics
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CBI 2000 Mosaic Power Spectrum
Mosaic Field Observations
• 3 fields totaling 40 deg^2
• Fields at d~0 a=2h, 14h, 20h
• ~125 nights of observing
• ~ 600,000 uv points covariance matrix 5000 x 5000
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CBI 2000 Mosaic Power Spectrum
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Cosmological Parameters
wk-h: 0.45 < h < 0.9, t > 10 Gyr
HST-h: h = 0.71 ± 0.076
LSS: constraints on s8 and G from 2dF, SDSS, etc.
SN: constraints from Type 1a SNae
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SZE Angular Power Spectrum
[Bond et al. 2002]
•Smooth Particle Hydrodynamics
(5123) [Wadsley et al. 2002]
•Moving Mesh Hydrodynamics
(5123) [Pen 1998]
•143 Mpc s8=1.0
Dawson et al. 2002
•200 Mpc s8=1.0
•200 Mpc s8=0.9
•400 Mpc s8=0.9
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Constraints on SZ “density”
• Combine CBI & BIMA (Dawson et al.) 30 GHz with
ACBAR 150 GHz (Goldstein et al.)
• Non-Gaussian scatter for SZE
– increased sample variance (factor ~3))
• Uncertainty in primary spectrum
– due to various parameters, marginalize
• Explained in Goldstein et al. (astro-ph/0212517)
• Use updated BIMA (Carlo Contaldi)
Courtesy Carlo Contaldi (CITA)
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SZE with CBI: z < 0.1 clusters
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New : Calibration from WMAP Jupiter
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Old uncertainty: 5%
2.7% high vs. WMAP Jupiter
New uncertainty: 1.3%
Ultimate goal: 0.5%
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New: CBI 2000+2001 Results
Future plans
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CBI 2000+2001 Noise Power
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CBI 2000+2001 and WMAP
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CBI 2000+2001, WMAP, ACBAR
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The CMB From NRAO HEMTs
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Post-WMAP Unification
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CBI + COBE
.
weak
2000
prior
2000+2001
0.12
k
-0.01
ns
2
2
b h

.
2000
+0.07
2000+2001
+0.09
-0.10
+0.01
+0.07
-0.12
-0.12
-0.08
-0.06
-0.07
0.09
+0.10
+0.11
+0.11
1.05
cdm h
weak prior + LSS
1.03
1.02
1.05
-0.08
-0.08
-0.07
-0.08
0.08
+0.05
+0.03
+0.026
0.17
0.11
0.12
0.10
-0.06
-0.03
-0.03
-0.021
0.015
+0.013
+0.014
+0.013
0.022
0.040
0.026
0.043
-0.009
-0.014
-0.010
-0.013
0.25
+0.15
+0.11
+0.08
0.40
0.62
-0.27
weak prior: t > 1010 yr
0.45 < h < 0.9
m > 0.1
0.64
-0.23
0.67
-0.14
-0.10
LSS prior: constraint on amplitude of s8 and
shape of Geff (Bond et al. Ap.J. 2003)
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weak prior + WMAP c
.
WMAP
WMAP+CBI
0.050
k
-0.063
ns
2
2
b h

-0.071
-0.028
-0.023
0.032
+0.022
0.975
cdm h
+0.064
0.962
-0.020
-0.013
0.015
+0.072
0.125
0.120
-0.0092
-0.0092
0.0012
+0.0010
0.0234
0.0231
-0.0008
-0.0005
0.243
+0.289
0.437
0.446
-0.075
weak prior: t > 1010 yr
0.45 < h < 0.9
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m2003
> 0.1
-0.059
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CBI Current & Future
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CBI Polarization
Noise Temperature (K)
Ka-band Receiver
• CBI instrumentation
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12
10
8
6
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4
2
0
– Use quarter-wave devices for linear to circular conversion
– Single amplifier per receiver: either R or L only per element
2000 Observations
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One antenna cross-polarized in 2000 (Cartwright thesis)
Only 12 cross-polarized baseline (cf. 66 parallel hand)
Original polarizers had 5%-15% leakage
Deep fields, upper limit ~8 mK
2002 Upgrade
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26 –
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Upgrade in 2002 using DASI polarizers (switchable)
Observing with 7R + 6L starting Sep 2002
Raster
28 scans for
30 mosaicing
32 and efficiency
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New TRW InP HEMTs Frequency
from NRAO
(GHz)
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Polarization Sensitivity
CBI is most sensitive at the peak of the polarization power spectrum
TE
The compact configuration
EE
Theoretical sensitivity ±1s of CBI in
450 hours (90 nights) on each of 3
mosaic fields 5 deg sq (no
differencing), close-packed
configuration.
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Polarization Interferometry
“Cross hands” sensitive to linear polarization (Stokes Q and U):
where the baseline parallactic angle is defined as:
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E and B modes
• A useful decomposition of the polarization signal is
into gradient and curl modes – E and B:
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CBI-Pol 2000 Cartwright thesis
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Pol 2003 – DASI & WMAP
Courtesy Wayne Hu – http://background.uchicago.edu
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CBI-Pol 2002-2004 Projections
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Conclusions from CBI Data
• Definitive measurement of diffusive damping scale
• Measurements of 3rd & 4th Acoustic Peaks
• At Low L  consistent with other experiments
• At High L (>2000)  indications of secondary anisotropy?
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Conclusions from CBI Data
• Definitive measurement of diffusive damping scale
• Measurements of 3rd & 4th Acoustic Peaks
• At Low L  consistent with other experiments
• At High L (>2000)  indications of secondary anisotropy?
Small Scale Power
• ~3 sigma above expected intrinsic anisotropy
• Not consistent with likely residual radio source populations
(more definitive characterization needed)
• Suggestive of secondary SZ anisotropy, although this would
imply sigma8 ~ 1
• Other possible foregrounds not ruled out at this point
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The CBI Collaboration
Caltech Team: Tony Readhead (Principal Investigator), John Cartwright, Alison Farmer, Russ
Keeney, Brian Mason, Steve Miller, Steve Padin (Project Scientist), Tim Pearson, Walter Schaal,
Martin Shepherd, Jonathan Sievers, Pat Udomprasert, John Yamasaki.
Operations in Chile: Pablo Altamirano, Ricardo Bustos, Cristobal Achermann, Tomislav Vucina,
Juan Pablo Jacob, José Cortes, Wilson Araya.
Collaborators: Dick Bond (CITA), Leonardo Bronfman (University of Chile), John Carlstrom
(University of Chicago), Simon Casassus (University of Chile), Carlo Contaldi (CITA), Nils
Halverson (University of California, Berkeley), Bill Holzapfel (University of California, Berkeley),
Marshall Joy (NASA's Marshall Space Flight Center), John Kovac (University of Chicago), Erik
Leitch (University of Chicago), Jorge May (University of Chile), Steven Myers (National Radio
Astronomy Observatory), Angel Otarola (European Southern Observatory), Ue-Li Pen (CITA),
Dmitry Pogosyan (University of Alberta), Simon Prunet (Institut d'Astrophysique de Paris), Clem
Pryke (University of Chicago).
The CBI Project is a collaboration between the California Institute of Technology, the Canadian
Institute for Theoretical Astrophysics, the National Radio Astronomy Observatory, the
University of Chicago, and the Universidad de Chile. The project has been supported by funds
from the National Science Foundation, the California Institute of Technology, Maxine and Ronald
Linde, Cecil and Sally Drinkward, Barbara and Stanley Rawn Jr., the Kavli Institute,and the
Canadian Institute for Advanced Research.
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