Transcript Présentation PowerPoint

The Virgo detector: status
and first experimental results
Nicolas Arnaud
NIKHEF
June 20th, 2003
Outline
• The quest for gravitational waves (GW): a long history
• Detection principle
 Interferometric detectors
• Description of the Virgo interferometer
 Optical scheme
 Main features of the instrument
 Foreseen sensitivity
• Experimental control of the Central Interferometer (CITF)
 CITF description and CITF commissioning goals
 Experimental results (spring 2001  summer 2002)
• Virgo versus the other GW interferometric detectors
 The LIGO interferometers (USA) + TAMA (Japan)
• Main GW sources and filtering techniques
Do gravitational waves exist?
• First «imagined» by Poincaré in 1905
été d'abord
conduitby
à supposer
que
la propagation de
• «J'ai
GW existence
predicted
Einstein in
1918
la gravitation n'est pas instantanée mais se fait à la vitesse
•de
A la
difficult
appearance
lumièrefirst
(…) Quand
nous parlerons donc de la position ou
laValidity
General
Relativity
linearization
?!
de
vitesseof
duthe
corps
attirant,
il s'agira
de cette position
ou
de cette
à l'instant
oùofl'onde
gravifique
est partie de
«GW vitesse
travel
mind »
Sir A.S. Eddington
• 50’s-60’s:
backatinthe
thespeed
footlights
ce corps (…)» [Italics of the author]
GW theoretical framework developped (Pirani & Isaacson)
• The breakthrough:
the binary pulsar PSR 1913+16 (1974)
Indirect evidence that GW exist
Hulse & Taylor (Nobel 1993)
[& Damour]
Yes they do!
20 years of measurement
GW main characteristics
• Perturbations of the Minkowski metric
• Quadrupolar emission
• Extremely weak!!!
Luminosity  G/c5  10-53 W-1
Ex: Jupiter radiates 5.3 kW as GW
during its orbital motion
 over 1010 years: EGW = 2  1021 J  Ekinetic  2  1035 J
No Hertz experiment possible!
Astrophysical sources required
• A good source of GW must be:
 asymetric
 compact (R ~ RSchwartzchild = 2GM/c2)
 relativistic
GW detectable effect
GW effect : differential modification of lengths
2
D
L
(
t
)
h(t) 
L
L
L + DL
h: dimensionless amplitude
h  1 / distance
The detector sensitivity volume should ultimately extend
beyond the Virgo cluster (~ 20 Mpc  65106 light years)
Two main categories of detectors:
• resonant bars
• giant interferometers, Earth-based or space-based
Virgo
LISA
A very large GW frequency domain
Frequency Range
GW ‘Probe’
• Extremely Low
Frequencies
10-18  10-15 Hz
CMB polarization
• Very Low Frequencies
10-9  10-7 Hz
Pulsar timing
• Low Frequencies
10-4  10-1 Hz
• High Frequencies
1  104 Hz
LISA
Earth-based detectors
Resonant bars or IFOs
Resonant bars
• First GW detectors:
Joe Weber’s pioneering work – see Phys. Rev. 117 360 (1960)
• Resonator:
GW
supraconducting
coupled with
deposit
cylindrical bar
a transducer
energy
inside
the bar
xm  M
m xM
Vibrations
modulate
DC voltage
• Network of bars working for years with high duty cycles
• Narrow-band sensitivities limited by noises difficult to beat
Interferometric detection
Suspended
Michelson
Interferometer
Mirrors used as
test masses
Incident GW
Sensitivity :
hsens 
Optical path
modification
Variation of the
power Pdet at the
IFO output port
1
Arm length  Power incident on BS
The Virgo optical scheme
Laser power: Pin = 20 W
Sensitivity h  1 / Pin
Gain : 3000  30  50 ~ 106
White
fringe
Laser
Detection
Photodiode
-17 / Hz
-23
-22
Sensitivity : hsens ~ 3 10-21
 To increase the arm length : 1 m  3 km
 To add Fabry-Perot cavities (Finesse = 50  Gain = 30)
 To add a recycling mirror (P = 1 kW on the Beam Splitter)
The Virgo SuperAttenuator
Length ~ 7 m; Mass ~ 1 ton
Structure in inverted pendulum
INFN
Pisa
-
g
k
1


fres 2π m
l
 fres ~ 30 mHz
Dual role:
• Passive seismic isolation
Seismic Attenuation:
~ 1014 at 10 Hz
• Mirror active control
only 0.4 N needed
for a 1 cm motion
Virgo foreseen sensitivity
Thermal
noise
Tail of the
0.6 Hz marionetta/
mirror resonance
Thermal
noise
Violin
mirrors
modes
Shot noise
«Seismic Wall»
Minimum ~ 3 10-23/
Hz between ~ 500 Hz et 1 kHz
The Virgo detector
Full Virgo
configuration
West Arm
Half-Arm Buildings
3 km
1.5 km
North Arm
3 km
1.5 km
Mode-Cleaner
144 m
Control
Building
Central Building
Virgo in numbers
• Arm length: 3 km
 6800 m3 in ultra-high vacuum (10-10 mbar)
• Very high quality mirrors:
 Diffusion < 5 ppm, absorption < 1 ppm
Fabry-Perot
 Reflectivity > 99.995%
 Radius of curvature 3450 m (4.5 mm sagitta) end mirrors
• Laser power: 20 W
• Seismic noise attenuation: > 1014 above 10 Hz
• Foreseen sensitivity range: 4 Hz  10 kHz
Best sensitivity ~ 3  10-23 / Hz around 1 kHz
• Control accuracy
 Length: down to 10-12 m
 Angular: from 10-6 to 10-9 radians
Status of Virgo
• Spring 2001-Summer 2002:
Successful commissioning of the central interferometer (CITF)
CITF: Virgo without the 3-km Fabry-Perot arms
But :
 Same suspensions
 Same control chain
 Ideal benchmark for the complete Virgo interferometer
• From autumn 2002: upgrade to Virgo
• March 2003: first beam in the 3-km arm
• The Full Virgo commissioning will start after summer
• First Physical Data: 2004 or a bit later…
Virgo central interferometer (CITF)
• CITF commissioning = 1rst step of Virgo commissioning
• Recycled and suspended Michelson Interferometer
• Uses the technology developped for the Virgo control system
• CITF commissioning goals:
 check the different component performances
 validate control algorithms
 test data management (acquisition, storage…)
«West» Mirror
Arm
lengths
~6m
Recycling
Mirror
«North»
Mirror
The CITF is not sensitive enough:
no hope to collect data with GW signal!!!
CITF and working point
Best sensitivity :
• Michelson on dark fringe  control arm asymmetry: l2-l1
• Recycling cavity resonant (maximize the stored power)
 control IFO mean length: l0 + (l1+l2)/2
Very narrow Working Point
In addition: residual low frequency motion of mirrors (0.6 Hz)
 CITF active controls needed (local and global)
Goal :
Longitudinal control
«Locking »
 Resonant cavities
dl ~ 10-10 – 10-12 m
Angular control
«Alignment »
 Aligned mirrors
dq ~10-9 – 10-7 rad
The steps of the Virgo control
Control aim: to go from an initial situation with
random mirror motions to the Virgo working point
• Decreasing the residual motion separately for each mirror
 Local controls
+ First alignment of mirrors
• Lock acquisition of the cavities
• Check working point control stability
• Switch on the angular control
 Automatic Alignment
Switching from
local controls
to
global controls
Fabry
Perot
cavity
Cavity Control
M1 (r1, t1)
L
M2 (r2, t2)
Characteristic quantity: the finesse F F   r1 r2
1 - r1 r2
• Linear around resonance
• Linear region width  1 / F
• Slope increasing with F
The higher F, the more
difficult the cavity control
A finesse of 400 (aligned CITF)
is high for a suspended cavity
Pound-Drever error signal
First control of the Michelson
Fringe Counting
Fringe interval
Global~ 0.5 mm
Control
Time (s)
AC Power
Error signal
Time (s)
DC Power
Dark fringe
Time (s)
Interferometer
power output
June 13th 2001
First control of the recycled CITF
A complex problem:
Stored
• Two lengths to be controlled
instead
of~one
• Pmax
5.8 W
Power
 coupled error signals
 Gain ~ 70
• Narrow resonance of the recycling(Pcavity
(high finesse)
laser ~ 80 mW)
IFO
• Limited force available to act on mirrors
output
• Error signal ~ to the electronic noise
outside
• Dark
fringeresonance
power
[weak laser power + Recycling mirrorless
reflectivity
«dark» = 98.5%]
 unperfect contrast
Recycling
Main issues:
correction • Large fluctuations of
• To select the right resonance
[trigger on the stored power]
the stored power:
West
• Simultaneous
acquisition of the 2 cavity
controls
 low
feedback gain
correction
• Fast
damping of the 0.6 Hz pendulumresonance
excited
misalignments
each time the locking attempt fails
December 16th 2001
CITF main steps
• 5 Engineering Runs
• 3 days duration (24h/24h)
• ~ 1 TB data collected
Channel
«Physics» Control Monitoring
/ Engineering Run
type
~ 5 MBytes/s
Data
2%
61 %
37 %
~ 160 TB/an
fraction
• The 2 first in Michelson configuration (9/01 and 12/01)
• The 3 others Recycled configuration (4/02, 5/02 and 7/02)
Engineering Run ER0 ER1 ER2 ER3 ER4
Duty Cycle
98% 85% 98% 96% 77%
All sources of control losses understood
 Improvements in progress
CITF sensitivity improvements
ER
Best
Sensitivity
m/Hz
June 2001  July 2002
Factor 103
improvement
@ 10 Hz
8 10-12
E0
(@ 500 Hz)
5 10-12
E1
(@ 500 Hz)
E2
10-14
(@ 1 kHz)
E3
5 10-15
(@ 1 kHz)
E4
10-16
(@ 1 kHz)
Room for
many more
Improvements
Virgo foreseen
sensitivity
Factor 105
improvement
@ 1 kHz
From the CITF to the full Virgo
• CITF commissioning completed
• Large improvements in sensitivity in only one year
 Gain in ‘experimental experience’  many upgrades for Virgo
 CITF  Virgo will provide ‘free’ sensitivity improvements:
• Arm length: 6 m  3 km  gain of a factor 500 in h
• Fabry-Perot cavities: factor 30 in addition
• Reduction of laser frequency noise
 In reality, such gains are unfortunately not automatic:
• Some noises do not depend on the laser optical path
• Noise hunting is a very long work
 Virgo scheme more complicated (4 lengths instead of 2)
 Control acquisition procedures  from CITF (under study)
 Virgo can benefit from the other detector experiences
Virgo versus other interferometers
10-7
June-August 2002
October-November 2002
10-12
LIGO
TAMA
10-20
1 Hz
10 kHz
10-7
Virgo CITF
10-20
10 Hz
10 kHz
• All sensitivities in m/Hz
 Comparable plots!
• Improvements still needed!
• Record sensitivity: Tama
10-18 m/Hz @ 1 kHz
10-20 July
1 Hz
2002
5 kHz
• @ 10 Hz, the CITF has the
best sensitivity: 10-13 m/Hz
One word about LISA
• Constellation of 3 satellites
• 3 semi-independent IFOs
• Optimal combinations to
maximize SNR or study noise
• Search periodical sources
• Expected lifetime: 5 years
•Approved by NASA/ESA
• To be launched in 2011
• Earth-based detectors
limited by seismic noise
below few Hz
• Strong sources certainly
exist in the mHz range
Seismic wall
Preparing the GW Data Analysis
• Activity parallel to the experimental work on detectors
 1 international conference / year (GWDAW)
• Large number of potential GW sources:
 compact binary coalescences (PSR 1913+16)
 black holes
 supernovae
 pulsars
 stochastic backgrounds
…
• The corresponding signals have very different features
 various data analysis techniques
Coincidence detections
Why ?
• Some detectors
will be working
in the future
now ACIGA
 LIGO : 4 km
 VIRGO : 3 km
 GEO : 600 m
 TAMA : 300 m
 ACIGA : 500 m
• Coincidence = only way to separate a GW (‘global’ in the
network) from transient noises in IFOs
• Coincidences may allow to locate the source position in sky
• Coïncidences with other emissions: g, n
Declination d
Interferometer angular response
• 2 maxima
GW perpendicular
to detector plane
• 4 minima
blind detector!
e.g. when the GW
comes along the
arm bissector
Right ascension a
Reduction of a factor ~ 2 in average of the amplitude
Example of the Virgo-LIGO network
• Spatial responses
 in a given direction
• Similarities between
the maps of the two
LIGO interferometers
• Complementarity
Virgo / LIGO
 Good coverage of
the whole sky
 Double or triple
coincidences
unlikely
Summary
• Many interferometers are currently under developpement
 Worldwide network in the future
 All instruments work already although they did
not prove yet there can fulfill their requirements
 Control of complex optical schemes with suspended mirrors
 All sensitivities need to be significally improved to
reach the amplitude of GW theoretical predictions
• Many different GW sources
 various data analysis methods in preparation
• In the two last years, the Virgo experiment became real
 The different parts of the experiment work well together
 Successful commissioning of the CITF
 2003: CITF  Full Virgo
 First ‘physically interesting’ data expected for 2004 !?!?!
GW: a never ending story
The future of gravitational astronomy looks bright.
1972
That the quest ultimately will succeed seems almost assured.
The only question is when, and with how much further effort.
1983
[I]nterferometers should detect the first
waves in 2001 or several years thereafter (…)
1995
Km-scale laser interferometers are now coming on-line, and it
seems very likely that they will detect mergers of compact
binaries within the next 7 years, and possibly much sooner.
2002
Kip S. Thorne
References about Virgo and GW
• Virgo web site: www.virgo.infn.it
• Virgo-LAL web site (burst sources):
www.lal.in2p3.fr/recherche/virgo
• Source review: C. Cutler - K.S. Thorne, gr-qc/0204090
• Some other GW experiment websites:
 LIGO: www.ligo.caltech.edu
 GEO: www.geo600.uni-hannover.de
 TAMA: www.tamago.mtk.nao.ac.jp/tama.html
 IGEC (bar network): igec.lnl.infn.it
 LISA: sci.esa.int/home/lisa
• Moriond 2003: moriond.in2p3.fr/J03
«Gravitational Waves and Experimental Gravity»
Recent status of all detectors: bars, IFOs and LISA
Detector noise characterization
Gaussian noise characterization: Power Spectrum Density (PSD)
S n (f)  2  FT  A n () 
FT: Fourier Transform
one-sided PSD (only positive frequencies)
with A n ()  lim
T
T/2
T/2

n(t) n(t) dt Autocorrelation function
T
• If the noise is dimensionless, the PSD unit is Hz-1
• RMS in the bandwidth [f1;f2]: RMS  f1 ; f2  
• s~n(f) 
Sn(f)
Detector
Sensitivity:
ff2 Sn(f) df
1
Amplitude Spectrum Density (unit 1/ Hz )
Log-log
scales
graph
Sn
or
s~n
Frequency (Hz)
Compact binary coalescences
Chirp signal:
amplitude and frequency
increase with time until
the final coalescence
The signal knowledge ends
before the coalescence
when approximations used
for the computation are
no more valid.
 large theoretical work
to go beyond this limit!
Example: PSR 1913+16
Waveform analytically estimated by developments in v/c
Coalescence expected in a few hundred million years
 Wiener filtering used for data analysis
Virgo will (?!?) be sensitive to the last minutes…
Optimal but computationally expensive
Impulsive sources (‘bursts’)
Examples:
• Merging phase of binaries
• Supernovae
• Black hole ringdowns
GW main characteristics:
• Poorly predicted waveforms
 model dependent
• Short duration (~ ms)
• Weak amplitudes
 Need to develop  filters :
 robust (efficient for a large class of signals)
 sub-optimal (/ Wiener filtering)
 online (first level of event selection)
Zwerger
/ Müller
examples of
simulated
supernova
GW signals
Pulsars
• GW signal: permanent, sinusoidal, possibly 2 harmonics
• Weak amplitude  detection limited to the galaxy
• Matched filtering-like algorithms using FFT periodograms
• Idea: follow the pulsar freq. on large timescales (~ months)
 compensation of frequency shifts: Doppler effect
due to Earth motion, spindown…
• Very large computing power needed (~ 1012 Tflops or more)
 Hierarchical methods are being developped  1 TFlop
 Need to define the better strategy:
 search only in the Galactic plane, area rich of pulsars
 uniform search in the sky not to miss close sources
 focus on known pulsars
• Permanent signal  coincident search in a single detector:
compare candidates selected in 2 different time periods
Stochastic backgrounds
• Described by an energy density per unit logarithmic
frequency normalized to the critical density of the universe:
2
3
π f 3S stoch (f)
H
0
with  c  8G
• Two main origins: Ω stoch (f) 
2G ρc
 Cosmological
Emission just after the Big Bang: ~10-44 s, T~1019 GeV
Detection  informations on the early universe
 Astrophysical
Incoherent superposition of GW of a given type emitted
by sources too weak to be detected separately.
• Detection requires correlations between 2 detectors
• After 1 year integration: h02 stoch  10-7 (1rst generation)
10-11 (2nd generation)
• Theoretical predictions: ~ 10-13  10-6
• Current best limit: stoch  60 @ 907 Hz [Explorer/Nautilus]