Present and Future. Nergis Mavalvala October 09, 2002

Transcription

1 Gravitational-wave Detection with Interferometers Present and Future Nergis Mavalvala October 09,

2 Interferometric Detectors Worldwide LIGO TAMA LISA LIGO VIRGO GEO 2

3 Global network of detectors LIGO GEO VIRGO TAMA LIGO AIGO LISA 3

4 Gravitational Waves General relativity predicts transverse spacetime distortions propagating at speed of light In TT gauge and weak field approximation, Einstein field equations wave equation Conservation laws Conservation of energy no monopole radiation Conservation of momentum no dipole radiation Lowest moment of field quadrupole (spin 2) Radiated by aspherical astrophysical objects Radiated by dark mass distributions black holes, dark matter 4

5 Astrophysics with GWs vs. E&M E&M Space as medium for field Accelerating charge incoherent superpositions of atoms, molecules Wavelength small compared to sources images Absorbed, scattered, dispersed by matter 10 MHz and up Detectors have small solid angle acceptance GW Spacetime itself Accelerating aspherical mass coherent motions of huge masses Wavelength large compared to sources no spatial resolution Very small interaction; matter is transparent 10 khz and down Detectors have large solid angle acceptance Very different information, mostly mutually exclusive Difficult to predict GW sources based on E&M observations 5

6 Astrophysical sources of GWs Coalescing compact binaries Classes of objects: NS-NS, NS-BH, BH-BH Physics regimes: Inspiral, merger, ringdown Other periodic sources Spinning neutron stars numerically hard problem Burst events Supernovae asymmetric collapse Stochastic background Primordial Big Bang (t = sec) Continuum of sources The Unexpected GWs neutrinosphotons now 6

7 Strength of GWs: e.g. Neutron Star Binary Gravitational wave amplitude (strain) h 2G 4π = I& h 4 µν c r µν GMR 4 c r For a binary neutron star pair 2 2 f 2 orb M kg M R M R f 20 km 400 Hz h ~10-21 r r m 7

8 GWs meet Interferometers Laser interferometer λ L L P IN L- L L+ L φ=0 L = h L φ= φ Suspend mirrors on pendulums free mass Optimal antenna length L ~ λ/4 ~ 10 5 meter 8

9 Practical Interferometer For more practical lengths (L ~ 1 km) fold interferometer to increase phase sensitivity φ = 2 k L N (2 k L); N ~ 100 N number of times the photons hit the mirror Light storage devices optical cavities Dark fringe operation lower shot noise φ=n L GW sensitivity on beamsplitter P increase power Power recycling Most of the light is reflected back toward the laser recycle light back into interferometer Price to pay: multiple resonant cavities whose lengths must be controlled to ~ 10-8 λ 9

10 Power-recycled Interferometer Optical resonance: requires test masses to be held in position to meter Locking the interferometer end test mass Light bounces back and forth along arms ~100 times 30 kw Light is recycled ~50 times 300 W input test mass Laser + optical field conditioning 6W single mode signal 10

11 WA LIGO 4 km 3030 km (±10 ms) 2 km LA 4 km 11

12 Initial LIGO Sensitivity Goal Strain sensitivity < 3x /Hz 1/2 at 200 Hz Displacement Noise Seismic motion Thermal Noise Radiation Pressure Sensing Noise Photon Shot Noise Residual Gas Facilities limits much lower 12

13 Limiting Noise Sources: Seismic Noise Motion of the earth few µm rms at low frequencies Passive seismic isolation stacks amplify at mechanical resonances but get f -2 isolation per stage above 10 Hz 13

14 Limiting Noise Sources: Thermal Noise Suspended mirror in equilibrium with 293 K heat bath k B T of energy per mode Fluctuation-dissipation theorem: Dissipative system will experience thermally driven fluctuations of its mechanical modes: ~ h ( f ) = π k B fl T Re( Z(f) is impedance (loss) Low mechanical loss (high Quality factor) Suspension no bends or kinks in pendulum wire Test mass no material defects in fused silica Z ( f )) FRICTION 14

15 Limiting Noise Sources: Quantum Noise Shot Noise Uncertainty in number of photons detected 1 hc λ h ( f ) = 2 L 8 F P Higher input power P bs need low optical losses (Tunable) interferometer response T ifo depends on light storage time of GW signal in the interferometer Radiation Pressure Noise Photons impart momentum to cavity mirrors Fluctuations in the number of photons h( f Lower input power, P bs ) = 2F ML bs T 2hP 3 π c ifo bs λ ( T τ 1 ifo s (, τ f f s 2, ) f ) Optimal input power for a chosen (fixed) T ifo 15

16 25 cm diameter, 10 kg fused silica optics Polished substrates Core Optics Micro-roughness < 10 ppm scatter Optical coatings < 2 ppm scatter < 1 ppm absorption Metrology Surface uniformity ~1 nm rms 16

17 17

18 Displacement Sensitivity (Preliminary S1, Sept. 2002) 18

19 Evolution of strain sensitivity Major improvements Front-end electronics (ADC) noise whiten signal Laser frequency noise commonmode servo Output electronics (DAC) noise whiten signal Sensing electronics noise increase light level 19

20 LIGO I Instrument Status: What works All three interferometers locked for hours at a time in power-recycled configuration G rec ~ Lock acquisition is done using sys id methods MTTL ~ 1 2 minute Optical parameters consistent with lab metrology mirror losses < 70 ppm Mechanical (internal) modes of test masses excited 10 4 < Q < 10 7 End-to-end simulation program Data analysis pipelines and algorithms 20

21 LIGO I Instrument Status: What doesn t (yet) Large factor to be gained in strain sensitivity Laser input power 1 5 W but all light power not detected (PD saturations) alignment control system Coupling to environment, e.g. Earthquakes Trains twice a day at Livingston Anthropogenic noise seismic pre-isolation system Electronics improvements 21

22 Engineering runs The Task Ahead Characterize and improve detector sensitivity and reliability Exercise data analysis system end-to-end pipeline Science runs Upper limits (E7 and beyond) Scientific searches (S1 was 08/23/02 to 09/07/02) Coincidence with GEO and TAMA Commissioning remaining subsystems Factor of 100 (above 400 Hz) to (at 40 Hz) improvement needed to reach design sensitivity LIGO II R&D already underway 22

23 The next-generation detector Advanced LIGO (aka LIGO II) Now being designed by the LIGO Scientific Collaboration Goal: Quantum-noise-limited interferometer Factor of ten increase in sensitivity Factor of 1000 in event rate. One day > entire 2-year initial data run Schedule: Begin installation: 2006 Begin data run:

24 A Quantum Limited Interferometer Facility limits Quantum Gravity gradients LIGO I LIGO II Residual gas (scattered light) Advanced LIGO Seismic noise Hz Thermal noise 1/15 Suspension thermal Optical noise 1/10 Seismic Test mass thermal Beyond Adv LIGO Thermal noise: cooling of test masses Quantum noise: quantum non-demolition 24

25 Seismic noise How will we get there? Active isolation system Mirror suspended as fourth (!!) stage of quadruple pendulum Thermal noise Suspension: fused quartz; ribbons Test mass: higher mechanical Q material, e.g. sapphire Quantum noise Input laser power: increase to ~200 W Optimize interferometer response, T ifo : signal recycling 25

26 Optimizing the optical response: Signal Tuning r(l).e iφ (l) Power Recycling Cavity forms compound output coupler with complex reflectivity. Peak response tuned by changing position of SRM l Signal Recycling Reflects GW photons back into interferometer to accrue more phase 26

27 Advance LIGO Sensitivity: Improved and Tunable Thorne 27

28 Implications for source detection NS-NS Inpiral Optimized detector response NS-BH Merger NS can be tidally disrupted by BH Frequency of onset of tidal disruption depends on its radius and equation of state broadband detector BH-BH binaries ~10 min 20 Mpc 300 Mpc ~3 sec Merger phase non-linear dynamics of highly curved space time broadband detector Supernovae Stellar core collapse neutron star birth If NS born with slow spin period (< 10 msec) hydrodynamic instabilities GWs 28

29 Spinning neutron stars Source detection Galactic pulsars: non-axisymmetry uncertain Low mass X-ray binaries: If accretion spin-up balanced by GW spindown, then X-ray luminosity GW strength Does accretion induce non-axisymmetry? Stochastic background Can cross-correlate detectors (but antenna separation between WA, LA, Europe dead band) Ω(f ~ 100 Hz) = 3 x 10-9 (standard inflation ) GW energy / closure energy Thorne (primordial nucleosynthesis d 10-5 ) (exotic string theories 10-5 ) Sco X-1 Signal strengths for 20 days of integration 29

30 Detection of candidate sources Thorne 30

31 When gravitational waves are detected Tests of general relativity Waves direct evidence for time-dependent metric Black hole signatures test of strong field gravity Polarization of the waves spin of graviton Propagation velocity mass of graviton Astrophysical processes Inner dynamics of processes hidden from EM astronomy Cores of supernovae Dynamics of neutron stars large scale nuclear matter The earliest moments of the Big Bang Planck epoch Astrophysics 31

32 New Instrument, New Field, the Unexpected 32