WMAP 1st yrBOOMERanG 98

(1)

Effetto delle sorgenti sulla misura

di anisotropia CMB

(2)

b (deg) b (deg) b (deg)b (deg)

l (deg) 90GHz

l (deg)

b (deg)

220GHz

WMAP 1st yrBOOMERanG 98

b (deg)

l (deg) 150GHz

41GHz _{l (deg)} 60GHz _{l (deg)} 94GHz _{l (deg)}

(3)

l (deg) 90GHz

l (deg)

b (deg)

220GHz

b (deg)

l (deg) 150GHz

41GHz _{l (deg)} 60GHz _{l (deg)} 94GHz _{l (deg)} PKS0537-441

BOOMERanG 98WMAP 1st yr

(4)

l (deg) 90GHz

l (deg)

b (deg)

220GHz

b (deg)

l (deg) 150GHz

41GHz _{l (deg)} 60GHz _{l (deg)} 94GHz _{l (deg)} PMNJ0519-4546

BOOMERanG 98WMAP 1st yr

(5)

l (deg) 90GHz

l (deg)

b (deg)

220GHz

b (deg)

l (deg) 150GHz

41GHz _{l (deg)} 60GHz _{l (deg)} 94GHz _{l (deg)} PKS0454-46

WMAP 1st yrBOOMERanG 98

(6)

10 100 100 1000 10000

30 200

CMB rms

PKS0537-441 PMNJ0519-4546 PKS0454-46

μ K

CMB

in a 20' beam

frequency (GHz)

55 . 2 '

20

430 100

/ ⁻

⎟⎠

⎜ ⎞

⎝

= ⎛

⎟⎟⎠

⎜⎜ ⎞

⎝

⎛ Ω

GHz K

F

CMB

ν μ

(7)

• There are additional AGNs lost in the confusion of the CMB fluctuations.

• The WOMBAT catalogue and tools predict quite well the flux observed for the 3 detected

AGN, and can be used to estimate the contamination due to unresolved AGNs.

• In the 3% of the sky mapped by B98 the contamination of the PS at 150 GHz is less than 0.3% at l=200, and less than 8% at l=600.

• This is reduced by 50% if the resolved sources (at 150 GHz) are removed, and by 80% if

are removed those resolved at 41 GHz.

0.0 0.5 1.0 1.5 2.0 2.5 3.0 1

10 100

Flux (Jy) @ 150 GHz

counts

WOMBAT catalog

http://astron.berkeley.edu/wombat/foregrounds/radio.html

(8)

0 200 400 600 800 1000 1200 1400 0

1000 2000 3000 4000 5000

6000

_CMB

all sources

resolved @150GHz removed resolved @40GHz removed

l(l+1)c

l

/2 π ( μ K

2

)

multipole l

150GHz

(9)

Ma quante altre sorgenti ci sono, con flussi piu’ deboli, e quindi confuse nella CMB e nel rumore

del rivelatore ?

(10)

Log S Log N(>S)

DIPENDE DALLA LOGN-LOGS DELLE

SORGENTI CONSIDERATE And am

ento Eu

clid eo

Evoluzione

(11)

Survey radio e X

• Survey radio e X di AGN permettono di misurare la logN vs. logS da usare (vedi astro-ph/0306206).

• Per calcolare la varianza del flusso totale I

prodotto da tutte le sorgenti nelle diverse direzioni si deve calcolare

• dove

dS dS S dN

I

S

∫

=

Δ

^max

min

2 2

) (

)

( N S

dS d dS

dS dN dS

S dN N

S

>

=

⇒

=

> ∫

^∞

Dimensioni di N: 1/sr

Dimensioni di S: W/m

²

/Hz o Jy

quindi Jy

²

/sr (1 Jy=10

^-26

W/m

²

/Hz )

(12)

Fluttuazioni di flusso

• Il risultato e’ (per frequenze da 40 a 100 GHz)

• A questo punto si tratta di trovare qual’e la

fluttuazione di Temperatura CMB che produce lo stesso segnale sul rivelatore di questa

fluttuazione di flusso.

• La fluttuazione di potenza dovuta a Δ I e’

• La stessa fluttuazione potrebbe derivare da Δ T tale che

sr / Jy 30

²

2

≈

ΔI

2 2

2

A I

W = Ω Δ Δ

2 2

2

A B

W = Ω Δ

Δ

(13)

Fluttuazioni di flusso

• Dove

2 2

2

⎟ ⎠

⎜ ⎞

⎝

⎛

∂ Δ ∂

Ω

= Δ

Ω

= Δ

Δ Ω

= Δ

⎪⎩ ⇒

⎪ ⎨

⎧

Δ Ω

= Δ

Δ Ω

= Δ

T T B

B I

B B I

A W

I A

W

kT x hc

e xe T

T B hc e

T T

B

x x x

σ

σ ₌

= −

⎥ ⎦

⎢ ⎤

⎣

⎡

−

∂

= ∂

⎟ ⎠

⎜ ⎞

⎝

⎛

∂

∂ ;

1 )

( 2 1

2 3

sr

Fluttuazioni di flusso

• Quindi

• E per i coefficienti dello spettro di potenza

• E quindi

2 2

1 )

( ⎟⎟

⎠

⎜⎜ ⎞

⎝

⎛ Δ −

Ω

= Δ

Ω

=

Δ

_x ^x

e xe T

T T B

B I

2 ,

,

1 )

( ⎟⎟

⎠

⎜⎜ ⎞

⎝

⎛ Ω −

=

_T _x ^x

I

e

xe T

T c B

c

_l _l

2 ,

,

1 )

( ⎟⎟

⎠

⎜⎜ ⎞

⎝

⎛ Ω −

=

x x I

T

e xe T

T B

c

_l

c

^l

Jy

²

/sr

Jy

²

/(sr

²

K

²

) sr

K

²

(15)

Fluttuazioni di flusso

• Possiamo valutare c

_l,I

sapendo che le fluttuazioni possoniane producono fluttuazioni uguali a tutti i multipoli. Quindi deve essere

• Dove e’ la trasformata in armoniche sferiche della risposta angolare: nel caso Gaussiano

• E quindi

• Si ottiene quindi

( )

_, ² ^,

( )

²

2

2 1

1 4 4 2

1

l l

l

l c B

B c

I ⁼ ∑ ⁺

_I

⁼

^I

∑ ⁺

Δ π π

2

B

l

) 2

1 (

2 _b

e

B

_l

=

⁻^l ^l⁺ ^σ

( ) ( )

₂

0 )

1

2 (

1 2 1

2 1

2

² ² ²

b

x

dx

e x e

B

^b ^b

σ

≈ =

+

=

+ ∑

⁻ ⁺ ^∞

∫

⁻

∑

^l ^l

l l

l

l l

2 2

,

4 I

c

_l _I

= πσ

_b

Δ

(16)

Fluttuazioni di flusso

• Inoltre

• E quindi

2 0

2

0

2 0

2 2

) (

2 2

d

b

e

d d

sen e

d RA

b

πσ θ

θ π

ϕ θ

θ θ

σ θ

=

≈

≅

= Ω

∫

∞ −

∞

2 2

,

2 2

2 ,

,

1 )

/ ( 2

1 )

2 (

4 1 )

(

⎟⎟ ⎠

⎜⎜ ⎞

⎝

⎛ Δ −

=

⇒

⎟⎟ ⎠

⎜⎜ ⎞

⎝

⎛

−

= Δ

⎟⎟ ⎠

⎜⎜ ⎞

⎝

⎛ Ω −

=

x x T

x x b

b

x x I

T

e xe T

T I B

c

e xe T

T B

I

e xe T

T B c c

l

l l

πσ

(17)

• 1 Jy/sr=10

^-26

W/m

²

/sr/Hz

• B(ν,T)[Jy/sr] = B(ν,T)[W/cm

²

/sr/cm

^-1

] / 3x10

^-20

• Inserendo nella si ottiene:

2 2

, 1

) / (

2 ⎟⎟⎠

⎜⎜ ⎞

⎝

⎛ Δ −

= _x ^x

T e

xe T

T I B

asymmetry (resulting in 10

⁹

γ for each baryon and no antibaryons)

• Thermalized in the primeval fireball by repeated scattering against free electrons, and

released when the universe cooled down enough to produce atoms and become transparent

(380000 yrs after the big bang) .

• Redshifted to microwave

frequencies (z

_CMB

=1100) and diluted in the subsequent 13.7 Gyrs of expansion of the Universe

γ

→ 2 + b

b

t

10

⁻⁶

s

10

¹³

s

10

¹⁷

s

visible NIR

MW

visible NIR

MW

T=3000K

T=3K

em now em

now

r z = = r

+ λ

1 λ

T >1GeV

B(ν)

Today: 410 γ/cm

³

b

b −

(19)

The spectrum: a proof of the primeval fireball

wavenumber σ ( ^cm

^-1

)

Mather et al. 1994

(20)

0 5 brightness temperature of the sky (K)

at 150 GHz

• The CMB dominates the sky brightness at mm wavelengths

• And is very much isotropic: the early universe was very homogeneous

• The most boring picture of the sky ever !

(21)

The Cosmic Microwave Background:

Observing directly the Early Universe

• Why

– A direct view of the primeval plasma

(380kyr after the big‐bang, and 13.7Gyr ago)

– An indirect view of the first split‐second – ultra‐high‐energy physics – A diffuse backlight illuminating the first structures

– A testbench for cosmology and fundamental physics

• How

– CMB observables and ultimate limits

– Dry & cold sites; near‐space and deep‐space missions for the CMB – Mm‐wave telescopes and ultrasensitive detector arrays

– Polarimeters – Spectrometers

• What has been achieved to‐date

– COBE‐FIRAS; BOOMERanG, DASI, et al.; WMAP; Planck – Large Telescopes for SZ and small scale anisotropy

• Much more to come

– Planck

– SZ & CMB spectroscopy

– Inflation & CMB Polarimetry

(22)

The anisotropy: an image of the primeval fireball

• We should be able to see causal horizons in the image of the CMB.

• 380000 yrs after the big bang, regions separated by more than 380000 light years are not in causal contact yet.

• We see these regions from a distance of approximately 13.7 Glyrs, and the universe expanded by a factor 1100 since then, so we expect a typical angular size of the causal horizon of the order of 1°.

• Regions which have never been in causal contact before have a very similar physical temperature. Why ? Paradox of horizons.

• Need deeper observations, enhancing the contrast of the image.

1

o

13.7Gyrs 1100 kyrs

380 × ≅

θ =

(23)

The anisotropy: an image of the primeval fireball

• Different physical effects, all related to the small density fluctuations δρ / ρ present 380000 yrs after the big bang (recombination) produce CMB Temperature fluctuations:

• Scales larger than the horizon are basically frozen in the pre‐recombination era. Flat power spectrum of δ T/T at large scales.

• Scales smaller than the horizon undergo acoustic

oscillations during the primeval fireball. Acoustic peaks in the power spectrum of δ T/T at sub‐degree scales.

c n c

T

T r r

⋅

− +

= v

4 1 3

1

2 γ

γ

ρ δϕ δρ

δ

Sachs-Wolfe (gravitational

redshift)

Photon density fluctuations

Doppler effect from velocity

fields

(24)

After recombination, density perturbation can grow and create the hierarchy of structures we see in the nearby Universe.

Before recombination

After recombination

T < 3000 K T > 3000 K

overdensity

Due to gravity, Δρ/ρ increases, and so does T

Pressure of photons

increases, resisting to the compression, and the perturbation bounces back

T is reduced enough that gravity wins again

Here photons are not tightly coupled to matter, and their pressure is not effective.

Perturbations can grow and form Galaxies.

t t

Density perturbations (Δρ/ρ) were oscillating in the primeval plasma (as a result of the opposite effects of gravity and photon pressure).

(25)

Size of sound horizon

time

Big-bang recombination Power Spectrum

multipole220450

1st peak 2nd peak

LSS

300000 ly

In the primeval plasma, photons/baryons density perturbations start to oscillate only when the sound horizon becomes larger than their linear size . Small wavelength perturbations oscillate faster than large ones.

R

R C

C

1st dip 2nd dip

The angle subtendeddependson the geometryof space size of perturbation

(wavelength/2)

300000 y 0 y

v v

v

v v

v

(26)

Typical size of causal horizon :

1

o

13.7Gyrs 1100 kyrs

380 × ≅

θ =

• We cannot predict the exact pattern of the temperature fluctuations, because our theory is statistical.

• However we can predict

quite accurately the correlation properties of the image.

• Gaussian fluctuations: all the information is encoded in the power spectrum.

(27)

Critical density universe

Ω>1

High density Universe

l

Critical density Universe Low density Universe

PS

l PS

200 200 200

0 0 0

(29)

Normal Matter

4%

Dark Matter

22%

Dark Energy

74%

Radiation

< 0.3%

The power spectrum depends on the composition of the universe through the physics of the oscillations and the evolution of the bkg.

=1-Ω_Λ-Ω_b

(30)

WMAP

Bennett et al. 2003 Hinshaw et al. 2006

BOOMERanG

de Bernardis et al. 2000 Masi et al. 2005

1

^o

Detailed Views of the Recombination Epoch

(z=1088, 13.7 Gyrs ago)

(31)

Keisler et al. 2011, astro-ph/1105.3182

(32)

The Cosmic Microwave Background:

Observing directly the Early Universe

• Why

– A direct view of the primeval plasma –

sensitive to the geometry and composition of the universe

An indirect view of the first split‐second – ultra‐high‐energy physics – A diffuse backlight illuminating the first structures

• How

• What has been achieved to‐date

– COBE‐FIRAS; BOOMERanG, DASI, et al.; WMAP;

Large Telescopes for SZ and small scale anisotropy

• Much more to come

– Planck

(33)

The linear polarization: testing physics of the primeval fireball and near the big bang

• CMB photons were last scattered 380000 yrs after the big bang.

• It was a Thomson scattering.

• For a given scattering center, any quadrupole

anisotropy in the incoming photons produces linear polarization in the scattered photons.

• Two possible origins for quadrupole anisotropy:

– Density perturbations present at recombination

– Gravitational waves produced during the inflation process, a split‐second after the big‐bang (E=10

¹⁶

GeV?)

• Different symmetry properties

(34)

-

+

-

+ ^x

y

- -

+

-

+

x y

- x

y

-10ppm +10ppm

= e

^-

at last scattering

(35)

Komatsu et al. 2010 – astro-ph/1001.4538

+ +

- - +

- -

+

Velocity field near density fluctuation

Resulting anisotropy seen by e-

Resulting polarization pattern

WMAP7 measured data (stacked)

E-modes : 3 μK (2002…)

(36)

Cosmic Inflation :

An exponential, superluminal expansion of space, happening a split‐second after the big bang,

due to a phase transition of the universe, at Energies of the order of 10¹⁶GeV (!)

(37)

Expansion vs Horizon

time size of the horizon

size of

the considered region According to the inflation

theory ….

…had been causally connected to the surrounding regions before inflation

380000 y

A region as large as the horizon when the CMB is released ….

(38)

time size of the horizon size of

the considered region

10^-36 s

no rm al

evo lut ion

In fl a ti o n : e x p o n e n ti a l

e x p a n si o n

(39)

time size of the horizon size of

the considered region

10^-36 s

Here the horizon contains all of the universe observable today

In fl a ti o n : e x p o n e n ti a l

e x p a n si o n

no rm al

evo lut ion

(40)

Quantum fluctuations (pre‐inflation)

Density fluctuations (in the primeval fireball)

Inflation provides an explanation for the origin

of the density fluctuations producing CMB anisotropy

(41)

• If inflation really happened:

9 It stretched geometry of space to nearly Euclidean

9 It produced a nearly scale invariant spectrum of gaussian density

fluctuations

9 It produced a stochastic background of gravitational waves: Primordial G.W.

The background is so faint that even LISA will not be able to measure it.

• Tensor perturbations also produce

quadrupole anisotropy. They generate irrotational (E-modes) and rotational (B-modes) components in the CMB polarization field.

• Since B-modes are not produced by scalar fluctuations, they represent a signature of inflation.

Quadrupole from P.G.W.

E-modes

B-modes

(42)

E-modes & B-modes

• From the measurements of the Stokes Parameters Q

and U of the linear polarization field we can recover both irrotational and rotational a

_lm

by means of

modified Legendre transforms:

( ) ⁿ ( ^a ^ia ) ^Y ^{( )} ⁿ

iU

Q

_m

m

B m E

m

r r

l l

l

l 2

,

)

( ^± ⁼ ∑ ^±

_±

( ) [ ( ) ( ) ( ) ( ) ]

∫

− +

−

− +

Ω

=

− +

+ Ω

=

n Y

n iU

Q n

Y n

iU Q

n W i d

a

n Y

n iU

Q n

Y n

iU Q

n W d

a

m m

B m

m m

E m

r r

r

r r

r

l l

l

l l

l

2 2

) (

) 2 (

1 ) (

) 2 (

1

E-modes produced by scalar and tensor perturbations

B-modes produced only by tensor perturbations

Spin-2 quantity Spin-2 basis

(43)

• The amplitude of this effect is very small, but depends on the Energy scale of inflation. In fact the amplitude of

tensor modes normalized to the scalar ones is:

• and

• There are theoretical arguments to expect that the

energy scale of inflation is close to the scale of GUT i.e.

around 10

¹⁶

GeV.

• The current upper limit on anisotropy at large scales gives T/S<0.5 (at 2σ)

• A competing effect is lensing of E‐modes, which is important at large multipoles.

GeV 10

7 .

3

¹⁶

4 / 4 1

/ 1

2 2 4

/ 1

≅ ×

⎟⎟ ⎠

⎜⎜ ⎞

⎝

≡ ⎛

⎟ ⎠

⎜ ⎞

⎝

⎛ V

C C S

T

Scalar

GW

Inflation potential

B-modes from P.G.W.

⎥ ⎥

⎦

⎤

⎢ ⎢

⎣

⎡

≅ × +

GeV 10

1 2 . 2 0

) 1 (

16 4 / 1 max

K V

c

^B

μ

π

^l

l

(44)

The Cosmic Microwave Background:

Observing directly the Early Universe

• Why

– A direct view of the primeval plasma (380kyr after the big‐bang,13.7Gyr ago)

• How

• What has been achieved to‐date

• Much more to come

– Planck

(45)

S-Z

cluster

01 . 500 0

5

2

≈ =

Δ =

keV keV c

m kT

e

ν

e

ν

10

4

01 . 0 01

.

0 × =

⁻

Δ ≈ Δ ≈

ν τ ν

T T

Sunyaev R., Zeldovich Y.B., 1972, Comm. Astrophys. Space Phys., 4, 173 Birkinshaw M., 1999, Physics Reports, 310, 97-195

Incoming CMB photons

• Inverse Compton Effect for CMB photons against charged particles in the hot gas of clusters

• Cluster optical depth: τ=nσ

l l

= a few Mpc = 10²⁵ cm n < 10^-3 cm^-3

σ = 6.65x10^-25 cm²

• So τ = nσ

l

< 0.01 : there is a 1% likelihood that a CMB photon crossing the cluster is scattered by an electron

• E_electron >> E_photon, so the electron transfers energy to the photon. To first order, the energy gain of the photon is

• The resulting CMB temperature anisotropy is

Scattered CMB photons Incoming CMB photons

(46)

Brightness

All photons increase their energy. The result is a distortion of the spectrum of the CMB in the direction of rich clusters

A decrement at low frequencies

( <217GHz )

An increment at high frequencies

( > 217GHz )

frequency

Sunyaev

Sunyaev - - Zeldovich Zeldovich Effect Effect

(47)

• Being produced by scatterings, the S‐Z signal amplitude does not depend on the distance (redshift) of the cluster

• Depends linearly on the density of the gas

• The X‐ray brightness, instead, decreases significantly with distance and gas density (depends on the density

squared)

The Sunyaev‐Zeldovich Effect

X‐ray S‐Z

(48)

Atmospheric transmission, pwv=0.5mm

ISD Δ emission, 18K 6 kJy/sr @ 150 GHz

Thermal SZ

Non-Thermal SZ

Kinematic SZ

Thermal SZ

(49)

The Cosmic Microwave Background:

Observing directly the Early Universe

• Why

– A direct view of the primeval plasma (380kyr after the big‐bang,13.7Gyr ago)

• How

• What has been achieved to‐date

• Much more to come

– Planck

SZ Effect

(50)

How

(51)

Orders of magnitude :

• The CMB has a 2.725K blackbody spectrum

• Most common wavelength for a CMB photon: 2 mm

• Typical energy of a CMB photon: 0.6meV

• Density of CMB photons: 410 γ/cm

³

• Typical flux of photons: 10

¹²

γ/cm

²

/sr = 10

^‐10

W/cm

²

/sr

• Many photons, with very low energy.

• Detection difficult :

– Thermal Detectors : Bolometers

– Coherent detectors : antenna, (downconversion), amplifier

– Quantum detectors : KIDs where photons break Cooper pairs

(52)

•Environment difficult :

–Ultra‐cold & dry sites on the earth, or space –Cryogenics for detectors and optical systems

60 120 180 240 300 360 420

Frequency (GHz)

polarization

(53)

90μK 3μK

2 mm PWV

0.5 mm PWV

41 km

photon noise from atmosphere and the CMB

typical CMB anisotropy (90μK) and polarization (3μK) signals photon noise from mirrors (ε=0.005, ε x T !)

300K

40K

4K 1.5K

CMB

Left scale

Right scale

Fluctuations of the background

(54)

90μK 3μK

2 mm PWV

0.5 mm PWV

41 km

300K

40K

4K 1.5K

CMB

Left scale

Right scale

(60)

Detectors for the CMB

• A 50‐years‐long struggle to get to photon noise limited performance…

• … and then to increase the mapping speed,

replicating single pixels in large arrays.

(61)

1900 1920 1940 1960 1980 2000 2020 2040 2060

10

²

10

⁷

10

¹²

10

¹⁷

Langley's bolometer

Golay Cell

Boyle and Rodgers bolometer F.J.Low's cryogenic bolometer

Composite bolometer

Composite bolometer at 0.3K

Spider web bolometer at 0.3K Spider web bolometer at 0.1K 1year

1day 1 hour

1 second

Development of thermal detectors for far IR and mm-waves

ti m e r equi red to m a k e a m eas ur em ent ( s ec onds )

year

Photon noise limit for the CMB

(62)

Spider-Web Bolometers

Absorber

Thermistor

Built by JPL Signal wire

2 mm

•The absorber is micro machined as a web of metallized Si

₃

N

₄

wires, 2 μm thick, with 0.1 mm pitch.

•This is a good absorber for mm-wave photons and

features a very low cross section for cosmic rays.

Also, the heat capacity is reduced by a large factor with respect to the solid absorber.

•NEP ~ 2 10

^-17

W/Hz

^0.5

is achieved @0.3K

•150μK

_CMB

in 1 s

•Mauskopf et al. Appl.Opt.

36, 765-771, (1997)

(63)

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

= Photon noise limit

Multi-moded

(64)

What has been

achieved to‐date

(65)

MAT Atacama

CBI Atacama

DASI South Pole

Mainly coherent detectors ground- based in high-cold-dry sites

VSA Tenerife ACBAR South Pole

(66)

BOOMERanG : launched from McMurdo (Antarctica) 1998, 2003

(67)

<<3x10^-6 7x10^-8srad

1’

<<3x10^-4 7x10^-6 srad

10’

<<0.01 2x10^-4 srad

1^o

<<1 2x10^-2 srad

10^o

<RA_sidelobes>

Ω_mainlobe FWHM

Going to L2 reduces the solid angle occupied by the Earth by a factor 2π/2x10

^-4

=31000, thus

relaxing by the same factor the required off-axis rejection.

1.5Mkm

900km L2

COBE WMAP,

Planck

No day-night changes up there … extreme stability

(68)

WMAP (lauched in 2001)

(69)

Hinshaw et al. 2006

(70)

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, can explore

the CMB with high angular resolution (arcmin)

(71)

(72)

Atacama Cosmology Telescope 6m diameter, 1 deg² FOV

5190 m osl

South Pole Telescope

10m diameter, 1 deg² FOV 2800 m osl

sr cm

100

²

≅ Ω A

APEX @ Atacama

12m diameter, 0.3 deg² FOV 5190 m osl

(73)

10K

10K 0.3K

1m

(74)

sr cm

100 ²

≅ Ω

A

Survey Telescopes :

0.4 5900

1.3 230

0.7 2300

2.1 145

1.1 1000

3.1 95

FWHM(’) N

mm

GHz

(75)

(76)

(77)

(78)

(79)

(80)

No evidence for B-modes yet !

(81)

Planck is a very ambitious

experiment.

It carries a complex CMB experiment (the state of the art, a few years ago) all the way to L2, improving the sensitivity wrt WMAP by at least a factor 10,

extending the frequency

coverage towards high frequencies by a factor about 10

(82)

ESA : Jan Tauber

HFI PI : Jean Loup Puget (Paris)

HFI IS : Jean Michel Lamarre (Paris) LFI PI : Reno Mandolesi (Bologna) LFI IS : Marco Bersanelli (Milano)

Almost 20 years of hard work of a

very large team, coordinated by:

(83)

HFI LFI

Scientific Laboratories

Satellite

+ subcontractors National Agencies

PI Puget PI Mandolesi

(84)

Ecliptic plane

1 ^o/day

Boresight

(85^o from spin axis)

Field of view rotates at 1 rpm

E M

L2

Observing strategy

The payload works in L2, to avoid the

emission of the Earth, of the Moon,

of the Sun

(85)

HFI LFI

P P P P P P P

30 44 70 100 143 217 353 545 857

(86)

14 / May/ 2009

(87)

Thermal performance :

Planck collaboration: astro-ph/1101:2023

(88)

Thermal performance :

(89)

Mission :

(90)

(91)

Thermal SZ

Non-Thermal SZ

Kinematic SZ

Thermal SZ

best ground-based photometers: 4 bands

(92)

A2319

As seen by Planck

(93)

All-sky Sunyaev-Zeldovich clusters

• Planck multiband observations of SZ clusters (ESZ) over the full sky: 189 high quality cluster candidates detected

• The clusters in the ESZ sample are mostly at moderate redshifts lying between z=0.01 and z=0.55, with 86% of them below z=0.3. The ESZ-cluster masses span over a decade from 0.9 to 15 × 10

¹⁴

M

_sol

, i.e. up to the highest masses.

Discovered by Planck

Clusters : Planck collaboration: astro-ph/1101:2024

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All-sky Sunyaev-Zeldovich clusters

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Filaments

• Half of the baryons known to be present in the universe (from nucleosynthesis estimates) are missing (i.e. have not been detected in emission, nor in absorbtion).

• A possible physical state of

baryons escaping detection in the radio, IR, visible, X rays domains, is a ionized medium with low density or warm temperature.

• In principle, this is detectable in the microwaves, because it

produces a SZ effect.

• The Planck survey, observing the whole sky, has also observed couples of galaxy clusters, and there is evidence for filaments of gas connecting two of the couples (astro‐ph/1208.5911). This is hot gas, probably heated by shocks, but the density is low, so X‐ray emission is also very low. It might be one step towards the solution of the missing baryons problem.

106 YROSAT cts/s X rays

SZ

(96)

Sky coverage of the SPT SZ survey (astro‐ph/1210.7231)

2500 sq. deg = 9Mpixels !

(97)

• In the first 720 sq.deg.

analyzed:

– 230 cluster candidates

– 158 with counterparts found in the follow‐up optical/IR

observations – 121 have S/N > 5

– 117 new discoveries (!) – Cluster masses estimated:

• Comparing to models:

– The number of clusters that formed over the history of the universe is sensitive to the mass of neutrinos and the influence of dark energy on the growth of cosmic structures.

150 GHz SPT‐SZ survey

astro.ph/1203.5775

(98)

Ω

_Λ

=0.7 Ω

_Λ

=0.0

Simulations: e.g. Da Silva et al. astro‐ph/0011187

(99)

SPT SZ survey (550 sq.deg.)

134 new clusters discovered with the SZ effect (confirmed by optical and IR follow ups).

Reichart et al. astro‐ph/1203.5775

(100)

Much more

to come

(101)

The field is extremely active and growing No way to list every experiment,

currently working or planned.

A very biased personal selection follows …

DISCLAIMER

(102)

• Photometric observations of the SZ can be significantly biased, when there are less

spectral channels than free parameters.

• Components, LOS through a rich cluster:

ThSZ

KSZ

CMB

ISD

NThSZ

p

_min

, Amp

T

_d

, τ

_d

….( β )

At least, 8

independent

parameters !

(103)

Thermal SZ

Non-Thermal SZ

Kinematic SZ

Thermal SZ

best ground-based photometers: 4 bands

(104)

A2319

As seen by Planck

(105)

The final solution: spectroscopic measurements of the SZ

• Requirements:

– Wide spectral coverage

(in principle 100 to 1000 GHz)

– Modest spectral resolution (λ/Δλ = 100 to 1000) – Differential input, high rejection of common mode

signal (CMB is common mode and is 2750000 μK, cluster signal is differential and can be as low as 10 μK).

– Imaging instrument

– Wide field of view to image the whole cluster and

have a clean reference area to compare

(106)

• In fig. 1 we show the OLIMPO balloon

payload (Masi et al.

2008), with solar

panels, ground shield and sun shield

removed.

• Note the tiltable 2.6m primary mirror and the lightweigth secondary.

• Pointing is obtained rotating the payload around an azimuth

pivot and changing the elevation of the inner frame, including the telescope, the FTS and the detector’s cryostat

• The total mass of the

payload is 1.5 tons.

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(108)

The instrument is based on a double Martin Pupplett

Interferometer configuration to avoid the loss of half of the signal.

A wedge mirror splits the sky image in two halves I_A and I_B, used as input signals for both inputs of the two FTS’s.

Olimpo Telescope

Olimpo Cryostat

⎟⎟ ⎠

⎜⎜ ⎞

⎝

⎛ +

= 0

) 2 / sin(

) 2 /

cos( δ

_y

δ

FTSII

B

x

i A

E

⎟⎟ ⎠

⎜⎜ ⎞

⎝

⎛

= +

) 2 / sin(

) 2 / cos(

0 δ

δ

_x

y FTSI

A i E B

outgoing fields :

(109)

FTS Cold Optics

Optical optimization has been performed using ZEMAX^TM

software, optimizing the optical quality in the full FOV of

OLIMPO.

The instrument was designed to fit the available room in between the primary mirror and the cryostat, a 75x75x30 cm³ box.

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Simulated OLIMPO measurement of a cluster l.o.s. with τ_th=0.005,

T_e=10 keV, τ_nonth=0.0001, v_pec=500 km/s,

I_dust=6kJy/sr@150GHz The data with the error bars are simulated

observations from a single pixel of the OLIMPO-FTS, for an integration time of 3 hours. The two lines through the data points represent the input theory (thin) and the best fit for the plotted data realization (thick).

The other thin lines represent thermal plus non-thermal SZE, and dust emission.

this shift is due to the peculiar velocity of the cluster

The high-frequency excess is due to a modest amount of dust

dust thermal

+ non-thermal