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Cosmic Microwave Background

3/4

Paolo de Bernardis and Silvia MasiDipartimento di Fisica, Universita’ La Sapienza, Roma

Nizza, 13/Sep/2012

Yesterday :• We have seen how the CMB was

originated and its importance in modern cosmology

• We have seen how the spectralbrightness of the CMB has beenmeasured (with a cryogenicMartin-Pupplett interferometer, COBE-FIRAS)

• We have seen that we expect a low level of anisotropy of the image of the CMB, withstatistical properties dependingon the angular scale, with 1°being the size of the horizon at recombination.

• The spectrum

KTec

hTB

CMB

x

725.21

2),(3

2

=−

=νν

mmTBTB 06.1),(),( max =⇒= λλνλν

GHzkThx

CMBCMB 56

νν≅=

)31.5(159

82.23

1

1maxmax

maxmaxmax

−

−

==

⇒=⇒=−

cmGHz

xxe x

σν

WienRJ

E = 1 meV

CAMB codehttp://camb.info/

λw

l = 200 θ=1o

• mm-wavedetectors0.5 – 5 mm(bolometers &)

• extreme detector sensitivity: μK(cryo, space)

• Quest forangularresolution : <10’(μW-telescopes)

How to measure the image of the early universe ?

• T=2.725K:Millimeterwavelengths(0.001 eV photons)

• Low brightnesscontrast : COBE-DMR measured10ppm@10° scale

• Sub-degree-sized hot and cold spots (fromthe projected size of the causal horizon, 1°in a flat universe)

CMB Detectors• Quantum detectors use the interaction of photons with matter to

convert the photon energy in an electrical charge:– Photomultipliers: the binding energy of e- in metals is of the order of 1 eV

λ < 1 μm– Intrinsic Photoconductors: the binding energy of e- in crystals is of the

order of 0.1 eV λ < 10 μm– Doped Photoconductors: the binding energy of e- in doped crystals can be

as low as 0.01 eV λ < 100 μm– Kinetic Inductance Detectors: the binding energy of e- in Cooper-pairs is

of the order of 0.001 eV λ < 1000 μm


convert the photon energy in an electrical charge:– Photomultipliers: the binding energy of e- in metals is of the order of 1 eV λ

< 1 μm– Intrinsic Photoconductors: the binding energy of e- in crystals is of the order of

0.1 eV λ < 10 μm– Doped Photoconductors: the binding energy of e- in doped crystals can be as

low as 0.01 eV λ < 100 μm– Kinetic Inductance Detectors: the binding energy of e- in Cooper-pairs is of

the order of 0.001 eV λ < 1000 μm• Thermal Detectors use the integrated energy of a large number of

photons which are absorbed and heat-up a temperature transducer. – Bolometers use thermistors as sensors, either low temperature semiconductors

or superconductors.



λ < 1 μm– Intrinsic Photoconductors: the binding energy of e- in crystals is of the order

of 0.1 eV λ < 10 μm– Doped Photoconductors: the binding energy of e- in doped crystals can be as



photons which are absorbed and heat-up a temperature transducer. – Bolometers use thermistors as sensors, either low temperature

semiconductors or superconductors. • Coherent Detectors convert the EM wave in a current in an antenna,

and amplfy the current, possibly either directly with sufficiently fast amplifiers, or down-converting it at lower frequency using non-linear components (diodes).









and amplfy the current, possibly either directly with sufficiently fast amplifiers, or down-converting it at lower frequency using non-linear components (diodes). Suitable for CMB photons


convert the photon energy in an electrical charge.– Photomultipliers: the binding energy of e- in metals is of the order of 1 eV







and amplfy the current, possibly either directly with sufficiently fast amplifiers, or down-converting it at lower frequency using non-linear components (diodes). Suitable for CMB photons

Bolometers currently feature the best performance:Are more sensitive than coherent detectors for f>90GHz,Can be replicated in large arrays at low cost, do notdissipate significant power, and have been developedlonger than KIDs.

History: early days• The first bolometers were developed for

astronomy, and allowed the first IR spectroscopy of an astronomical source– Samuel Pierpoint Langley in 1878 develops the

bolometer: a thin blackened platinum strip, sensitive enough to measure the heat of a cow from a distance of ¼ mile.

– The detector works because the resistance of the Pt strip changes when heated by the absorbedradiation.

– The detector is differential: 4 strips are placed in a Wheatstone bridge but only one is blackenedand exposed to incoming radiation. Common-mode effects are rejected by the bridge and tinyvariations of bolometer resistance can bemeasured.

• With his bolometer Langley is able tomeasure the IR spectrum of the sun, discovering atomic and molecular lines.

One generation ago• The revolution :

– 1961: Franck J. Low develops the first cryogenic Ge bolometer, boosting the sensitivity by orders of magnitude.

– 1960’s and ff. bolometers and semiconductorsdetectors with their telescopes are carried tospace using stratospheric balloons and rockets.

• Consequence:– First sky surveys @ λ 100 μm

– 1968 First IR ground basedlarge area sky survey (2 μm, from Mt. Wilson)

Cryogenics is needed, to freeze noise of fundamental origin

• Johnson noise standard deviation : T0.5

• Phonon noise standard deviation : T• Photon noise (blackbody) standard deviation : T2.5

In addition, reducing T :• Makes the heat capacity of the radiation absorber

smaller, so faster detectors• Makes the dependance R(T) steeper for semiconductor

thermistors• Allows the use of superconducting transition thermistors

Modern bolometers work at T=0.3K or T=0.1K

Few decades ago

• mm-wave bolometers– cooled at 1.5K or 0.3K – operating from space

• become sensitive enough to measure the finestdetails of the Cosmic Microwave Background.

• Breakthrough:– The composite bolometer (absorber and thermistor

separated and each optimized independently): N. Coron, P. Richards …

Bolometersand the CMB:

F. Melchiorri (high mountain, 1974), ….

P. Richards et al.(balloon, 1980) …

and then John Mather etal. (1992) with the FIRAS on the COBE satellite:

these microwaves haveexactly a blackbodyspectrum

J. Mather : Nobel Prize in Physics, 2006

Circa 1970Composite Bolometer(Coron, Richards …)

Circa 1980monoliticbolometer(Goddard, ..)

15 years ago

• The spider-web absorber isdeveloped–It minimizes the heat capacity of the

absorber–It minimizes the cross-section to

cosmic rays, while maintaining high cross-section for mm-waves

Spider-web bolometers

Made in JPL

BOOMERanG 1998 (0.3K), Archeops 2001 (0.1K), ….Planck-HFI (lanuched 2009)

Spider-Web Bolometers

Absorber

Thermistor

Built by JPL Signal wire

2 mm

•The absorber is micromachined as a web of metallized Si3N4 wires, 2 μm thick, with 0.1 mm pitch.

•This is a good absorber formm-wave photons and features a very low cross section for cosmic rays. Also, the heat capacity isreduced by a large factorwith respect to the solidabsorber.

•NEP ~ 2 10-17 W/Hz0.5 isachieved @0.3K

•150μKCMB in 1 s

•Mauskopf et al. Appl.Opt. 36, 765-771, (1997)

Measured performance of Planck HFI bolometers (0.1K)(Holmes et al., Appl. Optics, 47, 5997, 2008)

=Photonnoiselimit

Multi-moded

Sensitivity to CMB anisotropy• A map of CMB anisotropy is a sampled image

ΔTi =ΔT(li,bi) for i=1,Npix , where ΔT(li,bi) is the average of ΔT(l,b) over the pixel area, for the pixel centered in (li,bi).

• Knowing : – the instantaneous sensitivity (NET), – the instrument angular resolution θ, – the sky coverage of the survey Ω

• we can compute the standard error for the estimate of ΔTi of each pixel, for a given total observation time t.

• Assuming uniform coverage and square pixels with side θ, we have simply

tNET

tNET

pixT

Ω==Δ θ

σ

Sensitivity to CMB anisotropy• Numerical example: assume

• You get

• Per pixel, over 14400 pixels: a large dataset, with a S/N ratio per pixel of the order of 3.

Kt

NETt

NET

stsKNET

pixT

oo

μθ

σ

θ

μ

27

'1200'12002020'10

104.3days 5150

5

=Ω

==

×=×=Ω

=×==

=

Δ

Sensitivity to CMB anisotropy• An array of n detectors optimally used will simply

multiply by n the observation time available foreach pixel.

• So we get

• The use of a large array can give more that just animprovement of sqrt(n). For ground basedobservations, atmospheric noise can besignificantly reduced by exploiting the correlationsof the noise over different pixels.

ntNET

ntNET

pixT

1θ

σ Ω==Δ

Today: Arrays of bolometers• When a detector becomes limited by the

fluctuations of the radiation it is detecting, there is no point in further improving the detector.

• The only way to improve the experiment, isto replicate the same detector in an array, boosting the mapping speed of the instrument.

• Superconducting thermistors can beproduced in the same automated processproducing the spider-webs, so TESs are becoming the standard technology for thisfield.

EBEX Focal Plane

• Total of 1476 detectors• Maintained at 0.27 K• 3 frequency bands/focal plane

738 element array 141 element hexagon Single TESLee, UCB

3 mm

5 cm

• G=15-30 pWatt/K • NEP = 1.4e-17 (150 GHz)• NEQ = 156 μK*rt(sec) (150 GHz)• msec, 3=τ

150

150 150

150250

250

420

Slide: Hanany

Using waveguide technologies• In place of the spider-web absorber …• … one can collect the incoming radiation using a

planar antenna, and convey the radiation on a planar superconducting strip-line towards a planar resistor mounted on a thermally insulatedisland with the thermistor.

• In this way the bolometer is smaller, with less heatcapacity

• Moreover, planar band-defining filters and channelizers can be placed on the same chip, thusresulting in a much more compact cold focalplanes.

• Example of this technology (among severalothers) : the SPIDER and Polar-Bear focal planes

Foal plane of the SPIDER experimentCaltech-JPL (J. Bock)2048 pixelsPlanar beamformingIntegrated band-definition filters and dual polarization

Focal plane for the Polarbear experimentUniversity of Berkeley (A. Lee)637 pixelslenslets + cross-slot antennas beamformingIntegrated dual polarization bolometersChannelizers being prepared.

Angular resolution: Telescopes for the CMB

• Large dimensions of the optical system collecting CMB radiation are required fortwo reasons :– To mitigate the effects of diffraction and detect

small structures in the CMB sky– To limit far sidelobes and reject strong signals

from the ground and other powerful sky sources(the Galactic plane, planets etc.).

θ = λ/D @λ=2 mm D=1.4m to have θ=5’

Importance of low sidelobes• The power detected is the integral of the brightness

times the solid angle, weighted with the angularresponse of the telescope:

• Typical telescoperesponse RA(θ,φ)

Ω= ∫ dRABAW ),(),(4

ϕθϕθπ

RA(θ)θmain lobe

side lobes

Brightness from direction (θ,φ)Telescope response

in direction (θ,φ)

FWHM=λ/D

boresight

Importance of low sidelobes• In the case of CMB

observations, the detectedbrightness is the sum of the brightness from the sky(dominant for the solid anglesdirected towards the sky, in the main lobe) and the Brightness from the ground(dominant for the solid anglesdirected towards ground, in the sidelobes).

RA(θ)θmain lobe

side lobes

FWHM=λ/D

boresight

⎥⎥⎥

⎦

⎤

⎢⎢⎢

⎣

⎡Ω+Ω= ∫∫ dRABdRABAW

lobesside

Ground

lobemain

sky ),(),(),(),( ϕθϕθϕθϕθ

Importance of low sidelobes

⎥⎥⎥

⎦

⎤

⎢⎢⎢

⎣

⎡Ω+Ω= ∫∫ dRABdRABAW

lobesside

Ground

lobemain

sky ),(),(),(),( ϕθϕθϕθϕθ

⎥⎦

⎤⎢⎣

⎡Ω+Ω≈

lobesside

lobessideGround

lobemain

lobemainsky RABRABAW ),(),(),(),( ϕθϕθϕθϕθ

signal of interest disturbance signal

K3≈ srad1<< sradπ2≈K300≈

signal of interest >> disturbance signal requires

),(),( ϕθϕθlobesside

lobemain RARA >>>

1≈

⎥⎦

⎤⎢⎣

⎡Ω+Ω≈

lobesside

lobessideGround

lobemain

lobemainsky RABRABAW ),(),(),(),( ϕθϕθϕθϕθ

signal of interest disturbance signal

K3≈ srad1<< sradπ2≈K300≈

600

)(

),(),(

),(),(srad

BB

RARA lobemain

Ground

sky

lobesside

lobemain

lobemain

lobesside

Ω≈⎥

⎦

⎤⎢⎣

⎡

⎥⎥⎥

⎦

⎤

⎢⎢⎢

⎣

⎡

Ω

Ω<<

ϕθϕθ

ϕθϕθ

1≈

FWHM Ωmainlobe <RAsidelobes>

10o 2x10-2 srad <<4x10-5

1o 2x10-4 srad <<4x10-7

10’ 7x10-6 srad <<1x10-8

1’ 7x10-8 srad <<1x10-10 !!!

What isRA(θ,φ) ?

• The intensity isthe square of the field:

• Example: a 2m diameter mirrorused at 1 cm and at 1 mm

2

1 )(2⎥⎦

⎤⎢⎣

⎡=

Ω θθ

akakJI

ddI

o

0.0 0.1 0.2 0.3 0.4 0.50.0

0.2

0.4

0.6

0.8

1.0

a=1m λ=1 cm λ=1 mm

angu

lar r

espo

nse

off-axis angle θ (deg)

mainlobe

sidelobes

What isRA(θ,φ) ?


10-3 10-2 10-1 100 101 10210-1410-1310-1210-1110-1010-910-810-710-610-510-410-310-210-1100

a = 1 m λ=1 cm λ=1 mm

angu

lar r

espo

nse


mainlobe

sidelobes

2

1 )(2⎥⎦

⎤⎢⎣

⎡=

Ω θθ

akakJI

ddI

o

What isRA(θ,φ) ?


• The first zero isfor

10-3 10-2 10-1 100 101 10210-1410-1310-1210-1110-1010-910-810-710-610-510-410-310-210-1100

a = 1 m λ=1 cm λ=1 mm

angu

lar r

espo

nse

off-axis angle θ (deg)a2

22.110λθ =

θ10

2

1 )(2⎥⎦

⎤⎢⎣

⎡=

Ω θθ

akakJI

ddI

o

What isRA(θ,φ) ?


• The first zero isfor

• The FWHM issimilar

10-3 10-2 10-1 100 101 10210-1410-1310-1210-1110-1010-910-810-710-610-510-410-310-210-1100

a = 1 m λ=1 cm λ=1 mm

angu

lar r

espo

nse

off-axis angle θ (deg)a2

22.110λθ =

0.5 FWHM

θ10

2

1 )(2⎥⎦

⎤⎢⎣

⎡=

Ω θθ

akakJI

ddI

o

What isRA(θ,φ) ?


• The envelope of the off-axisresponse scalesas θ-3

approximatelystarting from 0.5 at the FWHM

10-3 10-2 10-1 100 101 10210-1410-1310-1210-1110-1010-910-810-710-610-510-410-310-210-1100

a = 1 m λ=1 cm λ=1 mm

angu

lar r

espo

nse


3)( −∝θθRA

2

1 )(2⎥⎦

⎤⎢⎣

⎡=

Ω θθ

akakJI

ddI

o

Low diffraction design• Real world angular responses are worse than the one

studied here.• Sharp edges are in general important sources of

diffraction, and must be avoided in low sidelobesdesign. Use smoothed edges.

• A trumpet has a slow transition to free space at the aperture to avoid diffraction of sound waves.

• The spider supporting the secondary mirror in a Cassegrain telescope is an important source of diffraction.

• Penzias and Wilson used an under-illuminated off-axis paraboloid, to get low sidelobes

10dB = a factor 10 in power

Other example of low sidelobes design: Planck

STRAY LIGHT

Main Spillover

Main Beam

Sub Spillover

Main Spillover

F. Villa, LFI

Main Beam

Near Sidelobes

Angle from boresight

Resp

onse

Far Sidelobes

F. Villa, LFI

107

FWHM Ωmainlobe <RAsidelobes>

10o 2x10-2 srad <<1

1o 2x10-4 srad <<0.01

10’ 7x10-6 srad <<3x10-4

1’ 7x10-8 srad <<3x10-6

Going to L2 reduces the solid angle occupied bythe Earth by a factor 2π/2x10-4=31000, thusrelaxing by the same factor the required off-axisrejection.

1.5Mkm

900km L2

COBEWMAP,Planck

Telescopes for the CMB• After the COBE-DMR results there was a clear need

for meter-sized telescopes for the CMB.• Working at high frequencies requires smaller mirrors

for the same resolution. • A 1 m mirror at 150 GHz provides 10’ resolution, at

15 GHz provides only 1.4° .• However atmospheric noise at high frequencies is

severe. • So, waiting for a new space mission, two classes of

experiments were developed:– Ground-based radiometers working at high altitude

mountain sites, at λ around 1 cm – Balloon-borne bolometric receivers working at λ around 1-

2 mm

Many CMB telescopes !!

In different environments:high mountain / antarcticastratospheric balloonssatellites

Finally …• To be sure that what is detected is really CMB, one

MUST perform simultaneous measurements at different frequencies, to check the Spectrum of the anisotropy.

• So, experiments must be either multiband or evenspectroscopic.

• For example our BOOMERanG experiment had 4 bands matching atmospheric windows at balloonaltitude: 90 GHz, 140 GHz, 220 GHz, 410 GHz, and redundant bolometers for each band.

• This was the most effective strategy to fightsystematic effects:

BOOMERanG (1998, 2003)

Il lancio

Systematics ARE there. • In the real world noise is not gaussian and we have

drifts, spikes, events of different kind in the raw data.• Detectors characteristics (responsivity, noise) can

change with time during the survey. • Moreover, low-level local emission can contaminate

the sky signal in a non gaussian way.• Evident features are easily identified and rejected.• Features smaller than the noise cannot be removed,

and contaminate the results.• The experiment needs to have internal redundancy in

order to make tests for the presence of systematics.

Systematics ARE there. • The experiment needs to have internal redundancy in

order to make tests for the presence of systematics.A. Several detectors at the same frequencyB. Several different frequencies

• The experimental conditions must be changed, tocheck the reliability of the resultC. Experiment different scan speedsD. Experiment different sidelobes conditionsE. Experiment different locations of sun, moon,

strong sources.F. Results must be compared to results of similar,

independent experiments. • Calibration should be carried out several times during

the survey

Test A:

• Compare independent channels at the samefrequency.

• Different bolometers have different noiseperformance.

• Two channels with similar performance are B150A and (B150A1+B150A2)/2

• Sum and difference maps:

[B15

0A+(

B15

0A1+

B15

0A2)

/2]/2

[B15

0A-(

B15

0A1+

B15

0A2)

/2]/2

Test B:• The spectral test shows that the structures present in

the maps are CMB anisotropies. In fact:• The maps at different frequencies are plotted in

thermodynamic temperature units for the CMB (mK) so that structures with the spectrum of the CMB will appear the same at all frequencies.

• Structures with the spectrum of the CMB are evident in the maps and have high S/N at 90, 150, 240 GHz. The dust monitor channel at 410 GHzshows no CMB and very little dust.

x14.8 x14.8

“1%” Region

Are these genuine CMB fluctuations ?

The rms fluctuationsΔTrms = {Σl (2l+1) cl wl /4π}1/2

are spectrally distributed as the derivative of a 2.73K blackbody. All otherastrophysical sources of confusion do not fit the data.

This means that the bulk of the observed fluctuations has a cosmological origin.

2-D and 3-D scatter plotsconfirm this conclusion

10 10010-25

10-24

10-23

10-22

10-21 BOOMERanG-LDB

(spinning)dust

(thermal)dust

Synchrotron

Free-free

CMB

rms

Brig

htne

ss fl

uctu

atio

ns (W

/m2 /s

r/Hz)

frequency (GHz)Astro-ph/0011469

90GHz150GHz

240GHz

scatter plots ofhigh latitude data

Astro-ph/0011469

Test C

• We have a powerful tool: data were taken at twodifferent scan speeds: 1 dps and 2 dps.

• At 2dps the sky signal is converted into anelectrical signal at twice the frequency, whileinstrument related effects (transfer function, 1/f noise, microphonic lines etc.) remain at the samefrequency.

• For the same detectors compare maps from data taken at 1dps and from data taken at 2 dps

1 dps map + 2 dps map

1 dps map - 2 dps map

Test F:

BOOMERanG vs. WMAP

WMAP (2002)

Wilkinson Microwave Anisotropy Probe

WMAP in L2 : sun, earth, moon are allwell behind the solar shield.

WMAPHinshaw et al. 2006astro-ph/0603451

BOOMERanGMasi et al. 2005astro-ph/0507509

1oDetailed Views of the Recombination Epoch(z=1088, 13.7 Gyrs ago)

WMAP 3 years23-94 GHz

BOOMERanG-98145 GHz

BOOMERanG-03145 GHz

The consistency of the maps from three independentexperiments, working at very different frequencies and with very different mesurement methods, is the best evidence that the faint structure observed•is not due to instrumental artifacts•has exactly the spectrum of CMB anisotropy, so it isnot due to foreground emission•The comparison also shows the extreme sensitivity of cryogenic bolometers operated at balloon altitude (the B03 map is the result of 5 days of observation)

Hinshaw et al. 2006

The last revolution …ten years ago

• Large arrays of bolometers (2002 +)– TES allow complete microfabrication of

bolometers : large arrays possible– e.g. Caltech/JPL, Berkeley, NIST, Goddard,

Bonn, Paris, Grenoble …– The mapping speed is boosted.

• Coupled to large (10m) telescopes (2009+), can explore the CMB with high angularresolution (arcmin)

Atacama Cosmology Telescope6m diameter, 1 deg2 FOV5190 m osl

South Pole Telescope10m diameter, 1 deg2 FOV2800 m osl

Keisler et al. 2011, astro-ph/1105.3182

Keisler et al. 2011, astro-ph/1105.3182

Observations Matching TheoryFor anAdiabatic Inflationary Universewith acoustic oscillations

Normal Matter4%

DarkMatter

22%

DarkEnergy

74%

Radiation< 0.3%

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