Effetto delle sorgenti sulla misura di anisotropia CMB
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) 60GHzl (deg) 94GHzl (deg)
b (deg) b (deg) b (deg)b (deg)
l (deg) 90GHz l (deg)
b (deg)
220GHz
b (deg)
l (deg) 150GHz
41GHz l (deg) 60GHzl (deg) 94GHzl (deg) PKS0537-441
BOOMERanG 98WMAP 1st yr b (deg) b (deg) b (deg)b (deg)
l (deg) 90GHz l (deg)
b (deg)
220GHz
b (deg)
l (deg) 150GHz
41GHz l (deg) 60GHzl (deg) 94GHzl (deg) PMNJ0519-4546
BOOMERanG 98WMAP 1st yr
b (deg) b (deg) b (deg)b (deg)
l (deg) 90GHz l (deg)
b (deg)
220GHz
b (deg)
l (deg) 150GHz
41GHz l (deg) 60GHzl (deg) 94GHzl (deg) PKS0454-46
WMAP 1st yrBOOMERanG 98
100 10
100 1000 10000
200 30
CMB rms
PKS0537-441 PMNJ0519-4546 PKS0454-46
μK
CMBin a 20 ' bea m
frequency (GHz)
55 . 2 '
20
430100
/ −
⎟⎠
⎜ ⎞
⎝
= ⎛
⎟⎟⎠
⎜⎜ ⎞
⎝
⎛ Ω
GHz K
F
CMB
ν μ
• 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
0 200 400 600 800 1000 1200 1400 0
1000 2000 3000 4000 5000
6000
CMBall sources
resolved @150GHz removed resolved @40GHz removed
l(l+1)cl/2
π
(μ
K2 )multipole l
150GHz
Ma quante altre sorgenti ci sono, con flussi piu’ deboli, e quindi confuse nella CMB e nel rumore
del rivelatore ?
Log S Log N(>S)
DIPENDE DALLA LOGN-LOGS DELLE SORGENTI CONSIDERATE Andam
ento Eucl ideo Evoluzione
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
S
∫
=
Δ
maxmin 2 2
) ( )
( N S
dS d dS dS dN
dS S dN N
S
>
=
⇒
=
> ∫
∞Dimensioni di N: 1/sr Dimensioni di S: W/m2/Hz o Jy quindi Jy2/sr (1 Jy=10-26W/m2/Hz )
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
22
≈
ΔI
2 2
2
A I
W = Ω Δ Δ
2 2 2
2
A B
W = Ω Δ
Δ
Fluttuazioni di flusso
• Dove
2 2 2
2
2 2
2 2 2 2
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
3 2
sr
Jy2/sr K2 Jy2/(sr2K2)
Fluttuazioni di flusso
• Quindi
• E per i coefficienti dello spettro di potenza
• E quindi
2 2
2 2
1 )
( ⎟⎟
⎠
⎜⎜ ⎞
⎝
⎛ Δ − Ω
= Δ Ω
=
Δ
x xe xe T
T T B B
I
2 ,
,
1
) ( ⎟⎟ ⎠
⎞
⎜⎜ ⎝
⎛ Ω −
=
T x xI
e
xe T
T c B c
l l2 , ,
1 )
( ⎟⎟
⎠
⎜⎜ ⎞
⎝
⎛ Ω −
=
x x I T
e xe T
T B
c
lc
lJy2/sr
Jy2/(sr2K2) sr
K2
Fluttuazioni di flusso
• Possiamo valutare c
l,Isapendo 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 ,( )
22
2 1
1 4 4 2
1
l l
l l l l
l
l c B
B c
I = ∑ +
I=
I∑ +
Δ π π
2
B
l) 2 1 (
2
e
bB
l=
−ll+ σ( ) ( )
20 ) 1 (
2
2 1 2 1
1
2
2 2 2b
x
dx
e x e
B
b bσ
σ
σ
≈ =
+
=
+
−∞ +
−
∫
∑
∑
lll l l
l l
2 2
,
4 I
c
lI= πσ
bΔ
Fluttuazioni di flusso
• Inoltre
• E quindi
2 0
2
0 2 0
2 2
) (
2 2
2 2
d b
e
d d sen e d RA
b
b
πσ θ θ π
ϕ θ θ θ
σ θ
σ θ
=
≈
≅
= Ω
= Ω
∫
∫
∫
∞ −
∞ −
∞
2 2
,
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 πσ
• 1 Jy/sr=10-26W/m2/sr/Hz
• B(ν,T)[Jy/sr] = B(ν,T)[W/cm2/sr/cm-1] / 3x10-20
• Inserendo nella , 2 2 si ottiene:
1 ) / (
2 ⎟⎟⎠
⎜⎜ ⎞
⎝
⎛ Δ −
= x x
T e
xe T
T I B cl
10 100 1000
10
110
210
310
441 GHz 60 GHz 94 GHz 143 GHz 217 GHz 340 GHz 540 GHz CMB
l(l+1)Cl/2π (μK2)
multipole l
What is the CMB:
An abundant background of photons filling the Universe.• Generated in the very early universe, less than 4 μs after the Big Bang from a small asymmetry (resulting in 109γ 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 (zCMB=1100) and diluted in the subsequent 13.7 Gyrs of expansion of the Universe
γ
→ 2 b + b
t
10−6s
1013s
1017s
visible NIR MW
visible NIR MW T=3000K
T=3K
em now em now
r z
= =
r+ λ
1 λ
T >1GeV B(ν)
B(ν)
Today: 410 γ/cm3
b b
−
The spectrum: a proof of the primeval fireball
wavenumber
σ (
cm-1)
Mather et al. 1994
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 !
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
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
o13.7Gyrs 1100 kyrs
380 × ≅
θ =
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
After recombination, density perturbationcangrowand 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) wereoscillatingin the primeval plasma (as a result of the opposite effects of gravity and photon pressure).
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
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 v
v
Typical size of causal horizon :
1
o13.7Gyrs 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.
Critical density universe
Ω>1
Ω<1 High density universe
Low density universe 1o
2o
0.5o
Horizon
Ω=1 14 Gly
LSS
HorizonHorizon
The power spectrum depends on the global geometry of the universe
Ω>1 Ω=1 Ω<1
2o 1o
0.5o High density Universe l Critical density Universe Low density Universe PS
l PS
l PS
200 200 200
0 0 0
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-ΩΛ-Ωb
WMAP Bennett et al. 2003 Hinshaw et al. 2006
BOOMERanG de Bernardis et al. 2000 Masi et al. 2005
1o
Detailed Views of the
Recombination Epoch
(z=1088, 13.7 Gyrs ago)
Keisler et al. 2011, astro-ph/1105.3182
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
– 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;
Large Telescopes for SZ and small scale anisotropy
• Much more to come – Planck
– SZ & CMB spectroscopy – Inflation & CMB Polarimetry
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=1016GeV?)
• Different symmetry properties
-
-
+
-
+
xy
-
-+
-
+
x y
- x
y
-10ppm +10ppm
= e-at last scattering
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…) 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 (!)
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 ….
time size of the horizon size of
the considered region
10-36s
normal evolution
In
flation: exponential expansion
time size of the horizon size of
the considered region
10-36s
Here the horizon contains all of the universe observable today In
flation: exponential expansion
normal evolution
Quantum fluctuations (pre‐inflation)
Density fluctuations (in the primeval fireball) Inflation provides an explanation for the origin of the density fluctuations producing CMB anisotropy
• 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
E-modes & B-modes
• From the measurements of the Stokes Parameters
Q
andU
of the linear polarization field we can recover both irrotational and rotationala
lmby means of modified Legendre transforms:( ) n ( a ia ) Y ( ) n
iU
Q
mm
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 r r
r r
l l
l
l l
l
2 2
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
• 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
16GeV.
• 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
164 / 4 1 / 1
2 2 4 / 1
≅ ×
⎟⎟ ⎠
⎜⎜ ⎞
⎝
≡ ⎛
⎟ ⎠
⎜ ⎞
⎝
⎛
VC 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
Bμ
π
ll l
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)
– 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;
Large Telescopes for SZ and small scale anisotropy
• Much more to come – Planck
– SZ & CMB spectroscopy – Inflation & CMB Polarimetry
S-Z
cluster
01 . 500 0
5
2
≈ =
Δ =
keV keV c m
kT
e
ν
eν
10
401 . 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=nσl l= a few Mpc = 1025cm n < 10-3cm-3 σ = 6.65x10-25cm2
• So τ = nσl< 0.01 : there is a 1% likelihood that a CMB photon crossing the cluster is scattered by an electron
• Eelectron>> Ephoton, 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
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
• 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
X‐ray S‐Z
X‐ray S‐Z
Atmospheric transmission, pwv=0.5mm
ISD Δ emission, 18K 6 kJy/sr @ 150 GHz
Thermal SZ
Non-Thermal SZ
Kinematic SZ
Thermal SZ
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)
– 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;
Large Telescopes for SZ and small scale anisotropy
• Much more to come – Planck
– SZ & CMB spectroscopy – Inflation & CMB Polarimetry
SZ Effect
How
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
12γ/cm
2/sr = 10
‐10W/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 •Environment difficult :
–Ultra‐cold & dry sites on the earth, or space –Cryogenics for detectors and optical systems
60 120 180Frequency (GHz)240 300 360 420
polarization
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
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
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
K.Ka.Q-bands
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
V-band
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
W-band
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
D-band
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
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.
1900 1920 1940 1960 1980 2000 2020 2040 2060
10
210
710
1210
17Langley's bolometer Golay Cell
Golay Cell
Boyle and Rodgers bolometer F.J.Low's cryogenic bolometer
Composite bolometer Composite bolometer at 0.3K
Spider web bolometer at 0.3KSpider web bolometer at 0.1K 1year
1day 1 hour
1 second
Development of thermal detectors for far IR and mm-waves
time required to make a measurement (seconds)
year
Photon noise limit for the CMB
Spider-Web Bolometers
Absorber
Thermistor
Built by JPL Signal wire
2 mm
•The absorber is micro machined as a web of metallized Si3N4wires, 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 10-17W/Hz0.5is achieved @0.3K
•150μKCMBin 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)
= Photon noise limit
Multi-moded
What has been achieved to‐date
MAT Atacama
CBI Atacama
DASI South Pole
Mainly coherent detectors ground- based in high-cold-dry sites
VSA Tenerife ACBAR South Pole
BOOMERanG : launched from McMurdo (Antarctica) 1998, 2003
<<3x10-6 7x10-8 srad 1’
<<3x10-4 7x10-6srad 10’
<<0.01 2x10-4srad 1o
<<1 2x10-2srad 10o
<RAsidelobes>
Ωmainlobe 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
WMAP (lauched in 2001)
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, can explore the CMB with high angular resolution (arcmin)
Atacama Cosmology Telescope 6m diameter, 1 deg2FOV 5190 m osl
South Pole Telescope 10m diameter, 1 deg2FOV 2800 m osl
sr cm 100
2≅ Ω A
APEX @ Atacama 12m diameter, 0.3 deg2FOV 5190 m osl
10K
10K 0.3K
1m
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
No evidence for B-modes yet !
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
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:
HFI LFI Scientific Laboratories
Satellite
+ subcontractors National Agencies
PI Puget PI Mandolesi
Ecliptic plane
1 o/day
Boresight (85ofrom 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
HFI LFI
P P P P P P P
30 44 70 100 143 217 353 545 857
14 / May/ 2009
Thermal performance :
Planck collaboration: astro-ph/1101:2023
Thermal performance :
Planck collaboration: astro-ph/1101:2023
Mission :
Planck collaboration: astro-ph/1101:2022
Atmospheric transmission, pwv=0.5mm
ISD Δ emission, 18K 6 kJy/sr @ 150 GHz
Thermal SZ
Non-Thermal SZ
Kinematic SZ
Thermal SZ
best ground-based photometers: 4 bands
A2319 As seen by Planck
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 × 1014Msol, i.e. up to the highest masses.
Discovered by Planck
Clusters : Planck collaboration: astro-ph/1101:2024
All-sky Sunyaev-Zeldovich clusters
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.
106YROSAT cts/s X rays
SZ
Sky coverage of the SPT SZ survey (astro‐ph/1210.7231)
2500 sq. deg = 9Mpixels !
• 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
Ω
Λ=0.7 Ω
Λ=0.0
Simulations: e.g. Da Silva et al. astro‐ph/0011187
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
Much more to come
The field is extremely active and growing No way to list every experiment,
currently working or planned.
A very biased personal selection follows …
DISCLAIMER
DISCLAIMER
• 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 pmin, Amp
Td,
τ
d….(β
)At least, 8 independent parameters !
Atmospheric transmission, pwv=0.5mm
ISD Δ emission, 18K 6 kJy/sr @ 150 GHz
Thermal SZ
Non-Thermal SZ
Kinematic SZ
Thermal SZ
best ground-based photometers: 4 bands
A2319 As seen by Planck
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
• 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.
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 IAand IB, used as input signals for both inputs of the two FTS’s.
Olimpo Telescope
Olimpo Cryostat
⎟⎟⎠
⎜⎜ ⎞
⎝
⎛ +
= 0
) 2 / sin(
) 2 / cos(δ y δ
x
FTSII B iA
E
⎟⎟⎠
⎜⎜ ⎞
⎝
⎛
= +
) 2 / sin(
) 2 / cos(
0 δ
δ x
y FTSI
A i E B
outgoing fields :
FTS Cold Optics Optical optimization
has been performed using ZEMAXTM 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 cm3box.
Simulated OLIMPO measurement of a cluster l.o.s. with τth=0.005, Te=10 keV, τnonth=0.0001, vpec=500 km/s, Idust=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
OLIMPO spectrometer
Ideal ground-based 4-bands photometer
• Antenna diameter: 10 m
• Range of wavelengths:
0.01 – 20 mm
• Bolometric sensitivity (λ0.3mm, 1h integration):
5x10-9Jy
• Interferometry sensitivity (λ0.5mm, 300s integration, 16GHz bw) : 10-4Jy
• Interferometer beam:
10-9arcsec Millimetron ASC Moscow ROSCOSMOS
3 hours of observations of a rich cluster with a DFTS on Millimetron Absolutely outstanding. USING A PHOTON NOISE LIMITED BOLOMETER IN THE COLD ENVIRONMENT OF L2 WITH 4K TELESCOPE
3h integration on the same LOS through a rich cluster
P. de Bernardis, et al., Astronomy and Astrophysics, 538, A86 (2012)
7.5 ’
Oppure ammassi speciali …
1ES0657-556
Isolating SZ DM (at 223 GHz)
The SZE from the hot gas disappears at x0,th(∼ 220-223 GHz) while the SZDMexpected at the locations of the two DM clumps remains negative and with an amplitude and spectrum which depend on Mχ.
M
χ= 20 GeV M
χ= 40 GeV M
χ= 80 GeV
[Colafrancesco, de Bernardis, Masi, Polenta & Ullio 2006]
Inflation
Inflation & polarization & polarization
Kinetic Inductance Detectors MUSIC Caltech NIKA Inst.Neel Grenoble FBK Trento + Sapienza + INFN
CEB Idea and development: Leonid Kuzmin & (Chalmers + NijniNovgorod)
Polarization
Polarization modulators modulators
Pisano, G., Ng, M.W., Haynes, V., Maffei, B., “A Broadband Metal-Mesh Half-Wave Plate for Millimetre Wave Linear Polarisation Rotation”, Submitted to PIER JEMWA (2012)
EBEX EBEX Focal Plane
• Total of 1476 detectors
• Maintained at 0.27 K
• 3 frequency bands/focal plane
738 element array 141 element hexagon Single TES Lee, UCB
3 mm
5 cm
• G=15-30 pWatt/K
• NEP = 1.4e-17 (150 GHz)
• NEQ = 156 μK*rt(sec) (150 GHz)
•
τ = 3
msec,150
150 150
150 250
250 420
Slide: Hanany
William Jones Princeton University
for the Spider Collaboration
The Path to CMBpol June 31, 2009 Suborbital Polarimeter for Inflation Dust and the Epoch of Reionization
Suborbital Polarimeter for Inflation Dust and the Epoch of Reionization
Spider: A Balloon Borne CMB Polarimeter
• Long duration (~30 day cryogenic hold time) balloon borne polarimeter
• Surveys 60% of the sky each day of the flight, with ~0.5 degree resolution
• Broad frequency coverage to aid in foreground separation
• Will extract nearly all the information from the CMB E-modes
• Will probe B-modes on scales where lensing does not dominate
• Technical Pathfinder: solutions appropriate for a space mission
Carbon Fiber Gondola
Attitude Control
• flywheel
• magnetometer
• rate gyros
• sun sensor
Flight Computers/ACS
• 1 TB for turnaround
• 5 TB for LDB Pointing Reconstruction
• 2 pointed cameras
• boresight camera
• rate gyros Six single freq. telescopes
30 day, 1850 lb, 4K / 1.4 K cryostat
The Large The Large Scale Polarization Scale Polarization Explorer Explorer
P. deP. de BernardisBernardis, , forforthe LSPE the LSPE collaborationcollaboration SWIPE STRIP
LSPE in a nutshell
• The Large-Scale Polarization Explorer is
– a spinning stratospheric balloon payload – flying long-duration, in the polar night– aiming at CMB polarization at large angular scales – using polarization modulators to achieve high stability
• Frequency coverage: 40 – 250 GHz (5 channels)
• Angular resolution: 1.5 – 2.3 deg FWHM
• Sky coverage: 20-25% of the sky per flight
• Combined sensitivity: 10 μ K arcmin per flight
P. de Bernardis Bologna 14 feb 2012
SWIPE
• The Short WavelengthInstrument for the PolarizationExplorer
• Uses overmoded bolometers, trading angular resolution for sensitivity
• Sensitivity of photon-noise limited bolometers vs # of modes:
3.2 3.3 2.5 NET Focal Plane (μK/sqrt(Hz))=
30 25 NET (μK/sqrt(Hz) ) = 15
1.6 1.9 2.4 FWHM (deg) =
83 58 37 N det =
1.4 2.1 λ (mm) 3.3
220 145 90 f (GHz)
40 25 15 N modes (geom) =
m 0.8 F = Instrument
m 0.4 D lens = Bolometric
0.25 eff = LSPE - SWIPE
Number of modes actually coupling to the bolometer absorber P. de Bernardis Bologna 14 feb 2012
13
The EPIC-IM Mission Concept and
US CMB Activities
Jamie Bock (JPL/Caltech)
Beyond CORE Workshop
Paris, 25 June 2012 13
3.5 m 13 m
4 m Note: Configurations not shown on same scale Planck based on flight sensitivity and mission duration
EPIC- Low Cost Intermediate Mission 4 K Option Comprehensive Science Science Inflationary B-mode polarization
only
Inflationary B-modes, E-modes to cosmic variance, gravitational lensing to cosmic limits, neutrino mass, dark energy, Galactic astronomy
Inflationary B-modes, E-modes to cosmic variance, gravitational lensing, neutrino mass, dark energy, Galactic astronomy
Speed 140 Plancks 1000 Plancks 70 Plancks
Detectors 2400 11,000 (TES bolometer or MKID) 1500
Aperture Six 30 cm refractors 1.4 m Crossed Dragone telescope 3 m Gregorian Dragone Bands 30 – 300 GHz 30 – 300 GHz + 500 & 850 GHz 30 – 300 GHz
Cooling LHe cryostat + ADR 4 K Cryo-cooler + ADR TBD
Mass 1320 kg CBE 1670 kg CBE 3500 kg CBE
Publication ArXiv 0805.4207 (192 pages) ArXiv 0906.1188 (157 pages) ArXiv 0805.4207 (192 pages)
Cost $660M (FY07) $920M (FY09) No cost assessed
The EPIC-IM Design
13
Descoped Focal Plane for 30 K Telescope
‘4 K Telescope’ Option ‘30 K Telescope’ Option
• Focal plane reduced by 2.3x in mass
• Detectors become larger for 30 K telescope due to edge spillover
• Total detectors reduced from 11094 to 2022
COrE: www.core-mission.org ESA-M3 (2020)
There is still so much to discover about the early universe …
… and important technology improvements are making it possible !