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GW150914: The Advanced LIGO Detectors in the Era of First Discoveries
B. P. Abbott
et al.
(LIGO Scientific Collaboration and Virgo Collaboration)
(February 12, 2016)
Following a major upgrade, the two advanced detectors of the Laser Interferometer Gravitational-
wave Observatory (LIGO) held their first observation run between September 2015 and January
2016. With a strain sensitivity of 10
23
/
Hz at 100 Hz, the product of observable volume and mea-
surement time exceeded that of all previous runs within the first 16 days of coincident observation.
On September 14th, 2015 the Advanced LIGO detectors observed a transient gravitational-wave
signal determined to be the coalescence of two black holes [1], launching the era of gravitational-
wave astronomy. The event, GW150914, was observed with a combined signal-to-noise ratio of 24 in
coincidence by the two detectors. Here we present the main features of the detectors that enabled
this observation. At full sensitivity, the Advanced LIGO detectors are designed to deliver another
factor of three improvement in the signal-to-noise ratio for binary black hole systems similar in
masses to GW150914.
PACS numbers: 04.80.Nn, 95.55.Ym, 95.75.Kk, 07.60.Ly
Introduction
— On September 14th, 2015, both Ad-
vanced LIGO detectors in the USA, H1 in Hanford,
Washington and L1 in Livingston, Lousiana, made the
first direct measurement of gravitational waves [1]. The
event, GW150914, was determined to be the merger
of two black holes, with masses of 36
M
and 29
M
,
into a black hole of approximately 62
M
[2]. 3.0 so-
lar masses of energy (
'
5
.
4
×
10
47
J) was radiated in
gravitational waves. The gravitational waves from this
event, which occurred at a distance of
'
410 Mpc
'
1
.
3
×
10
9
light years, changed the separation between the
test masses by
'
4
×
10
18
m, about one 200th of a proton
radius.
The Advanced LIGO detectors,
multi-kilometer
Michelson-based interferometers [3], came online in
September 2015, after a major upgrade targeting a factor
of 10 sensitivity improvement over initial detectors [4, 5].
While not yet at design sensitivity during their first ob-
servation run, they have already exceeded the strain sen-
sitivity of the initial detectors across the entire frequency
band, significantly surpassing the past discovery poten-
tial [6–10]. This paper describes the Advanced LIGO
detectors, as well as their current and final design sensi-
tivity, at the inception of gravitational-wave astronomy.
Astrophysical Reach
— In general relativity, a gravita-
tional wave far from the source can be approximated as
a time-dependent perturbation of the space-time metric,
expressed as a pair of dimensionless strain polarizations,
h
+
and
h
×
[12]. An interferometric gravitational-wave
detector acts as a transducer to convert these space-time
perturbations into a measurable signal [13]. The inter-
ferometer mirrors act as ‘freely falling’ test masses. Ad-
vanced LIGO measures linear differential displacement
along the arms which is proportional to the gravitational-
wave strain amplitude. We define the differential dis-
placement as ∆
L
=
δL
x
δL
y
, where
L
x
=
L
y
=
L
are
the lengths of two orthogonal interferometer arms. The
gravitational-wave strain and the interferometer displace-
ment are related through the simple equation ∆
L
=
hL
,
where
h
is a linear combination of
h
+
and
h
×
.
The tiny displacements induced by astrophysical
events demand that the interferometer mirrors be free
from environmental disturbances and require a highly
sensitive interferometric transducer—designed to be lim-
ited only by disturbances arising from fundamental
physics consideration.
Since the interferometer re-
sponse to displacement, or equivalently gravitational-
wave strain, is frequency dependent, it is necessary to
represent the limiting detector noises as functions of fre-
quency normalized by the interferometer response.
The left panel of Figure 1 shows the amplitude spectral
density of the total strain noise in units of strain per
Hz during the first observation run (O1 run) and, for
comparison, during the final science run of the initial
LIGO detectors (S6 run). In the detectors’ most sensitive
frequency band between 100 Hz and 300 Hz, the O1 strain
noise is 3 to 4 times lower than achieved in the S6 run.
At 50 Hz, the improvement is nearly a factor of 100.
The right panel of Figure 1 shows the single detec-
tor signal-to-noise ratio (SNR) for an optimally oriented
compact binary system consisting of two 30
M
black
holes as a function of redshift
z
, and for different in-
terferometer configurations. The observed strain ampli-
tude is largest for a source whose orbital plane is parallel
to the detector’s plane and is located straight above or
below; we refer to such a source as optimally oriented.
The SNR is computed in the frequency domain [14] us-
ing standard cosmology [15] and phenomenological wave-
forms which account for inspiral, merger and ringdown,
but not spins [16].
A Michelson interferometer lacks good directional sen-
sitivity to gravitational waves. The antenna pattern cov-
ers approximately half the sky, both above and beneath
the plane of the detector. Moreover, the antenna patterns
arXiv:1602.03838v1 [gr-qc] 11 Feb 2016
2
10
100
1000
10
-
24
10
-
23
10
-
22
10
-
21
Frequency
(
Hz
)
Strain Noise
(
1
/
Hz
)
S6 run
O1 run
Adv. LIGO design
Future upgrades
0.03
0.1
0.3
1
3
10
30
100
300
1000
Redshift
SNR
GW150914
GW150914
(
S6
)
S6 Threshold
(
SNR
=
8
)
30
M
Black Hole Binaries
FIG. 1. The left plot shows the strain sensitivity during the first observation run (O1) of the Advanced LIGO detectors and
during the last science run (S6) of the initial LIGO detectors. The O1 strain noise curve is shown for H1 (dark red) and L1
(light red); the two detectors have similar performance. The Advanced LIGO design sensitivity as well as a possible future
upgrade [11] are shown to highlight the discovery potential in the coming years. The right plot shows the single detector
signal-to-noise ratio (SNR) under optimal orientation as function of redshift
z
—for two merging black holes with mass 30
M
each. GW150914 was not optimally orientated and was detected with a single detector SNR of 13 to 20 at
z
= 0
.
09; this event
would not have been seen in S6.
of the two LIGO detectors are aligned to maximize the
coincident detection of gravitational-wave signals, con-
strained to the 10 ms inter-site propagation time. The
coincidence constraint substantially rejects non-Gaussian
noise and vetoes local transients.
The observed strain amplitude is inversely propor-
tional to the luminosity distance. For small redshifts,
z <
1, the observable volume, and thus the detection
rate, grows as the cube of the detector sensitivity. The
number of detected events is expected to scale with the
product of observing volume and observing time. Be-
tween September 12 and October 20 the H1 and L1 de-
tectors had a duty cycle of 70% and 55%, respectively,
while the observing time in coincidence was 48%. Af-
ter data quality processing [17], 16 days of data were
analyzed around GW150914, resulting in a time-volume
product of 0.1 Gpc
3
yr for binary black hole systems with
masses similar to GW150914 [18].
The Displacement Measurement
— The current gener-
ation of advanced detectors uses two pairs of test masses
as coordinate reference points to precisely measure the
distortion of the space-time between them. A pair of in-
put and end test masses is located in each of the two arms
of a Michelson laser interferometer, as shown in Figure 2.
The Advanced LIGO test masses are very pure and ho-
mogeneous fused silica mirrors of 34 cm diameter, 20 cm
thickness and 40 kg mass.
It is critical that the test masses be free from sources of
displacement noise, such as environmental disturbances
from seismic noise, or thermally driven motion. These
noise sources are most relevant at frequencies below
100 Hz, while shot noise of the optical readout is dom-
inant at high frequency. Figure 3 shows the measured
displacement noise of Advanced LIGO during the first
observing run, together with the major individual con-
tributions, as discussed below.
To minimize ground vibrations, the test masses are
suspended by multi-stage pendulums [19], thus acting
as free masses well above the pendulum resonance fre-
quency of 0.4 Hz. Monolithic fused silica fibers [20] are
incorporated at the bottom stage to minimize suspension
thermal noise [21], which limits the useful frequencies to
10 Hz and above. The Advanced LIGO test masses re-
quire about 10 orders of magnitude suppression of ground
motion above 10 Hz. The multi-stage pendulum system
attenuates the ground motion by seven orders of magni-
tude. It is mounted on an actively controlled seismic iso-
lation platform which provides three orders of magnitude
of isolation of its own [22, 23]. Moreover, these platforms
are used to reduce the very large displacements produced
by tidal motion and microseismic activity. Tidal forces
can produce displacements up to several 100
μ
m over a
multi-kilometer baseline on time scales of hours. The
dominant microseismic activity is driven by ocean waves.
The resulting ground motion can be as large as several
μ
m at frequencies around 0.15 Hz—even far inland.
The entire test mass assembly including the suspension
system and part of the seismic isolation system resides
inside an ultra-high vacuum system, with pressures typi-
cally below 1
μ
Pa over the 10
,
000 m
3
volume, to prevent
acoustic shorting of the seismic isolation systems and to
minimize Rayleigh scattering in the optical readout.
The test masses are also susceptible to changes in the
local gravitational field caused by changing mass distri-
butions in their vicinity. While not limiting presently, at
design sensitivity this time-dependent Newtonian noise
3
Input test mass
End test mass
L
x
=
4
km
Michelson perpendicular arm
Laser
Output port
L
y
=
4
km
Suspension
System
Stage
1
Stage
2
Stage
3
1
.
6
m
Stage
4
Metal masses
Fused silica
masses
Steel wire
Fused silica fibers
Electrostatic
actuator
FIG. 2. Interferometer configuration and test mass setup.
Each arm of the Michelson interferometer includes two sus-
pended test masses. The two test masses are placed 4 km
apart and form an optical resonator with a gain of 300. The
suspension system is shown on the right, each test mass is at
the bottom of a quadruple pendulum. It provides high isola-
tion above the resonance frequencies which range from 0.4 Hz
to 13 Hz. The test mass is attached to the penultimate mass
through fused silica fibers providing a high mechanical qual-
ity factor which lowers the thermal noise. The other stages
use steel wire. The attachment point to the seismic isolation
system as well as stages 1 and 2 implement cantilever springs
for vertical isolation. Each test mass is accompanied by its
own reaction chain to minimize actuation noise. Coil actu-
ators are mounted to the upper stages of the reaction chain
and an electrostatic actuator is implemented at the bottom
stage. Shown on the left are the other optics of the Michelson
interferometer with the beamsplitter and the perpendicular
arm. The two optics at the interferometer input and output
port comprise the coupled resonator system which amplifies
the response of the optical transducer.
source possibly becomes relevant below 20 Hz, and might
require active cancellation [24, 25].
Thermally driven motion is another important source
of displacement noise. It includes the Brownian motion of
the suspension system [26] as well as the test masses [27],
and mechanical loss in the mirror optical coatings [28].
The mirror coatings, a dielectric multilayer of silica and
titania-doped tantala [29, 30], were developed to provide
the required high reflectivity while minimizing coating
thermal noise [31–33]; it limits the design sensitivity in
the central frequency band [3].
The predicted levels for seismic, thermal and Newto-
nian noise sources are summarized in Figure 3 and com-
pared to the total measured displacement noise. They are
currently not limiting the sensitivity due to the presence
of other technical noise sources, as detailed in Ref. [34].
Quantum noise in the interferometer arises from the
discrete nature of photons and their Poisson-distributed
arrival rate [35–37]. The momentum transfer of individ-
ual photons hitting a test mass gives rise to radiation
pressure noise. Quantum radiation pressure noise scales
as 1
/mf
2
, where
m
is the mass of the mirror and
f
the
frequency, and therefore it is most significant at lower
frequencies.
Photon shot noise arises from statistical fluctuations
in the photon arrival time at the interferometer output,
and it is a fundamental limit of the transducer in sensing
changes of the differential arm length. The importance of
shot noise decreases as the inverse square-root of the laser
power circulating in the interferometer arms. During the
first observing run, Advanced LIGO was operating with
100 kW of circulating laser power. The corresponding
quantum noise curve, comprising both low frequency ra-
diation pressure noise and high frequency shot noise, is
shown in Figure 3; it is limiting at frequencies above
100 Hz. In the upcoming years, we plan to increase the
circulating laser power up to 750 kW, and thus reducing
the shot noise contribution.
Coincident detection between the two LIGO obser-
vatories is used to reject transient environmental dis-
turbances. Both observatory sites deploy seismometers,
accelerometers, microphones, magnetometers, radio re-
ceivers, weather sensors, AC-power line monitors, and
a cosmic ray detector for vetoes and characterization of
couplings [38].
Interferometric Transducer
— The Advanced LIGO
detector uses a modified Michelson laser interferometer to
translate strain into an optical phase shift [3]. Similar to
an electromagnetic receiver, the optimal antenna length
for a gravitational-wave detector is a quarter wavelength.
For a gravitational wave at 100 Hz this is 750 km. The
Advanced LIGO interferometer arms are 4 km long and
employ an optical resonator between the input and end
test masses that multiplies the physical length by the ef-
fective number of round-trips of the light. However, the
physical length cannot be arbitrarily short, because test
mass displacement noises are multiplied by the same fac-
tor.
The output port of the Michelson interferometer is
held at an offset from a dark fringe, resulting in a small
amount of light leaving the output port [39]. A differen-
tial optical phase shift will then decrease or increase the
amount of light, depending which interferometer arm is
momentarily stretched or squeezed by a passing gravita-
tional wave. This light signal is measured by a photode-
tector, digitized and calibrated [40], before being sent to
the analysis pipelines [41, 42].
The calibration factor that converts detected laser light
power to mirror displacement is measured by applying
a known force to a test mass [43]. An auxiliary 1047-
nm wavelength laser is reflected off the end test mass
and modulated in intensity to generate a varying force.
The response of the optical transducer is measured by
sweeping the modulation frequency through the entire
detection band. It is also tracked by a set of fixed fre-
quency lines. This way, the calibrated strain readout is
4
Measured
Quantum
Thermal
Seismic
Newtonian
Other
DOF
10
100
1000
10
-
20
10
-
19
10
-
18
10
-
17
10
-
16
Frequency
(
Hz
)
Displacement
(
m
/
Hz
)
FIG. 3. The displacement sensitivity of the Advanced LIGO
detector in Hanford during the first observation run O1; the
Livingston detector has a similar sensitivity, as shown in Fig-
ure 1. The sum of all known noise sources accounts for most
of the observed noise with the exception of the frequency band
between 20 Hz and 100 Hz. This will be the focus of future
commissioning to full sensitivity. The quantum noise includes
both shot noise and radiation pressure noise. Thermal noise
includes terms due to the suspensions, the test masses and
the coatings. Seismic noise is the ground displacement atten-
uated through the seismic isolation system and the suspen-
sions. Cross couplings from the auto-alignment system and
from the auxiliary lengths are combined into the trace labelled
“other DOF” (degrees-of-freedom). Newtonian gravitational
noise is estimated from density perturbations due to surface
ground motion. The strong line features are due to the violin
modes of the suspension wires, the roll and bounce modes of
the suspensions, the AC power line and its harmonics, and
the calibration lines. Not shown are numerous noise sources
that don’t contribute significantly—such as laser frequency,
intensity and beam jitter noise, sensor and actuation noise,
and Rayleigh scattering by the residual gas [34].
computed in real-time with less than 10% uncertainty in
amplitude. The overall variability of the detector’s sen-
sitivity was about
±
10%.
The main light source is a pre-stabilized 1064-nm wave-
length Nd:YAG laser. It is followed by a high power am-
plifier stage, capable of generating a maximum output
power of 180 W [44]. During the first observation run,
only 20 W were injected into the interferometer. A tri-
angular optical resonator of 32.9 m round-trip length is
placed between the laser source and the interferometer to
reject higher order transverse optical modes and to stabi-
lize the laser frequency further [45]. At the output port,
a bow-tie optical resonator of 1.3 m round-trip length is
used to reject unwanted frequency components on the
light. Optical curvature mismatch of the interferometer
mirrors is caused by manufacturing imperfections and by
thermal lensing due to heating from the main laser beam.
A thermal compensation system provides active correc-
tion by means of ring heaters arranged around the test
masses and a set of CO
2
lasers for central heating [46].
The Advanced LIGO detector uses coupled optical res-
onators to maximize the sensitivity of the interferometric
transducer. These optical resonators enhance the light
power circulating in each arm while simultaneously opti-
mizing the effective antenna length and the gravitational-
wave signal bandwidth [47–51]. As the interferometer is
held near a dark fringe, most of the light is reflected
back to the laser source. Adding a partially transmissive
mirror at the interferometer input forms an optical res-
onator, leading to a power gain of 35 to 40 at the beam-
splitter. The optical resonator in the interferometer arms
enhances the circulating power by another factor of 300.
Thus, 20 W of laser power entering the interferometer
results in nearly 100 kW circulating in each arm. A par-
tially reflective mirror is also placed at the output port to
enhance the signal extraction and to increase the detec-
tor bandwidth. The resulting differential pole frequency
or detector bandwidth is
'
335 Hz (H1) and
'
390 Hz
(L1) [34].
All of these coupled optical resonators require servo
controls to be brought and held on resonance [52]. The
lengths of the optical resonators in the interferometer
arms are stabilized to less than 100 fm, whereas the
lengths of the other coupled resonators are kept within
1 to 10 pm [34]. Similarly, the interferometer test masses
are aligned within tens of nanoradians relative to the op-
tical axis for optimal performance. The noise arising from
sensing and control of these extra degrees-of-freedom are
combined together in the curve labeled “other DOF”
in Figure 3.
The Pound-Drever-Hall reflection lock-
ing technique is used to sense the auxiliary longitudi-
nal degrees-of-freedom [53, 54], while an interferometric
wavefront sensing scheme is deployed for the alignment
system [55, 56]. Digital servo systems are used to feed
control signals back to actuators which steer the rela-
tive longitudinal positions and orientations of the inter-
ferometer mirrors. To prevent reintroducing ground mo-
tion onto the test masses, electrostatic actuators [57] are
mounted to a second quadruple pendulum known as the
reaction chain. Only test masses use reaction chains; all
other interferometer mirrors use coil actuators mounted
on a rigid structure surrounding the suspensions.
Servo controls are also necessary for damping the
plethora of normal modes of the pendular suspensions
and for stabilizing the seismic isolation system to an in-
ertial reference frame. Moreover, at high laser power,
optical springs introduce angular instabilities due to pho-
ton radiation pressure-induced torques acting on the mir-
rors [58, 59], while the mirror acoustic modes introduce
parametric instabilities [60, 61]. At the current laser
power only one acoustic mode is unstable which can be
tuned away by the ring heaters. Together with thermal
heating, angular optical springs and multiple paramet-
ric instabilities are the main challenges that need to be
overcome to increase the circulating laser power; both
will require active damping for stable operations.
5
Overall, more than 300 digital control loops with band-
widths spanning from sub-Hz to hundreds of kHz are
employed to keep each Advanced LIGO interferometer
operating optimally during observation. The digital con-
trols computers also serve as the data acquisition system
that continuously writes on the order of 10
5
channels of
time series data to disk, at a rate of
'
12 MB/s. It is
synchronized to GPS to better than 10
μ
s [40]. A state-
based automation controller provides hands-free running
during operations.
Outlook
— The global gravitational-wave network will
be significantly enhanced in the upcoming years. In
2016 Advanced LIGO will be joined by Advanced Virgo,
the 3 km detector located near Pisa, Italy [62]. The
Japanese KAGRA interferometer [63] and a possible
third LIGO detector in India [64] will provide a global
network that allows for improved parameter estimation
and sky localization [65]. Achieving design sensitivity
with the network of current detectors will define earth-
bound gravitational-wave astrophysics in the near future.
Looking further ahead, we can envision current technolo-
gies leading to a factor of two improvement over the
Advanced LIGO design sensitivity [11], so that events
such as GW150914 could be detected with SNRs up to
200. More dramatic improvements will require significant
technology development and new facilities.
Acknowledgement
— The authors gratefully acknowl-
edge the support of the United States National Science
Foundation (NSF) for the construction and operation of
the LIGO Laboratory and Advanced LIGO as well as
the Science and Technology Facilities Council (STFC)
of the United Kingdom, the Max-Planck-Society (MPS),
and the State of Niedersachsen/Germany for support of
the construction of Advanced LIGO and construction
and operation of the GEO600 detector. Additional sup-
port for Advanced LIGO was provided by the Australian
Research Council. The authors gratefully acknowledge
the Italian Istituto Nazionale di Fisica Nucleare (INFN),
the French Centre National de la Recherche Scientifique
(CNRS) and the Foundation for Fundamental Research
on Matter supported by the Netherlands Organisation
for Scientific Research, for the construction and opera-
tion of the Virgo detector and the creation and support
of the EGO consortium. The authors also gratefully ac-
knowledge research support from these agencies as well
as by the Council of Scientific and Industrial Research
of India, Department of Science and Technology, India,
Science & Engineering Research Board (SERB), India,
Ministry of Human Resource Development, India, the
Spanish Ministerio de Econom ́ıa y Competitividad, the
Conselleria d’Economia i Competitivitat and Conselleria
d’Educaci ́o, Cultura i Universitats of the Govern de les
Illes Balears, the National Science Centre of Poland, the
European Commission, the Royal Society, the Scottish
Funding Council, the Scottish Universities Physics Al-
liance, the Hungarian Scientific Research Fund (OTKA),
the Lyon Institute of Origins (LIO), the National Re-
search Foundation of Korea, Industry Canada and the
Province of Ontario through the Ministry of Economic
Development and Innovation, the Natural Science and
Engineering Research Council Canada, Canadian Insti-
tute for Advanced Research, the Brazilian Ministry of
Science, Technology, and Innovation, Russian Founda-
tion for Basic Research, the Leverhulme Trust, the Re-
search Corporation, Ministry of Science and Technology
(MOST), Taiwan and the Kavli Foundation. The au-
thors gratefully acknowledge the support of the NSF,
STFC, MPS, INFN, CNRS and the State of Niedersach-
sen/Germany for provision of computational resources.
This document has been assigned the LIGO Labora-
tory document number LIGO-P1500237.
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Authors
B. P. Abbott,
1
R. Abbott,
1
T. D. Abbott,
2
M. R. Abernathy,
1
F. Acernese,
3
,
4
K. Ackley,
5
C. Adams,
6
T. Adams,
7
P. Addesso,
3
R. X. Adhikari,
1
V. B. Adya,
8
C. Affeldt,
8
M. Agathos,
9
K. Agatsuma,
9
N. Aggarwal,
10
O. D. Aguiar,
11
L. Aiello,
12
,
13
A. Ain,
14
P. Ajith,
15
B. Allen,
8
,
16
,
17
A. Allocca,
18
,
19
P. A. Altin,
20
S. B. Anderson,
1
W. G. Anderson,
16
K. Arai,
1
M. C. Araya,
1
C. C. Arceneaux,
21
J. S. Areeda,
22
N. Arnaud,
23
K. G. Arun,
24
S. Ascenzi,
25
,
13
G. Ashton,
26
M. Ast,
27
S. M. Aston,
6
P. Astone,
28
P. Aufmuth,
8
C. Aulbert,
8
S. Babak,
29
P. Bacon,
30
M. K. M. Bader,
9
P. T. Baker,
31
F. Baldaccini,
32
,
33
G. Ballardin,
34
S. W. Ballmer,
35
J. C. Barayoga,
1
S. E. Barclay,
36
B. C. Barish,
1
D. Barker,
37
F. Barone,
3
,
4
B. Barr,
36
L. Barsotti,
10
M. Barsuglia,
30
D. Barta,
38
J. Bartlett,
37
I. Bartos,
39
R. Bassiri,
40
A. Basti,
18
,
19
J. C. Batch,
37
C. Baune,
8
V. Bavigadda,
34
M. Bazzan,
41
,
42
B. Behnke,
29
M. Bejger,
43
C. Belczynski,
44
A. S. Bell,
36
C. J. Bell,
36
B. K. Berger,
1
J. Bergman,
37
G. Bergmann,
8
C. P. L. Berry,
45
D. Bersanetti,
46
,
47
A. Bertolini,
9
J. Betzwieser,
6
S. Bhagwat,
35
R. Bhandare,
48
I. A. Bilenko,
49
G. Billingsley,
1
J. Birch,
6
R. Birney,
50
S. Biscans,
10
A. Bisht,
8
,
17
M. Bitossi,
34
C. Biwer,
35
M. A. Bizouard,
23
J. K. Blackburn,
1
C. D. Blair,
51
D. G. Blair,
51
R. M. Blair,
37
S. Bloemen,
52
O. Bock,
8
T. P. Bodiya,
10
M. Boer,
53
G. Bogaert,
53
C. Bogan,
8
A. Bohe,
29
P. Bojtos,
54
C. Bond,
45
F. Bondu,
55
R. Bonnand,
7
B. A. Boom,
9
R. Bork,
1
V. Boschi,
18
,
19
S. Bose,
56
,
14
Y. Bouffanais,
30
A. Bozzi,
34
C. Bradaschia,
19
P. R. Brady,
16
V. B. Braginsky,
49
M. Branchesi,
57
,
58
J. E. Brau,
59
T. Briant,
60
A. Brillet,
53
M. Brinkmann,
8
V. Brisson,
23
P. Brockill,
16
A. F. Brooks,
1
D. A. Brown,
35
D. D. Brown,
45
N. M. Brown,
10
C. C. Buchanan,
2
A. Buikema,
10
T. Bulik,
44
H. J. Bulten,
61
,
9
A. Buonanno,
29
,
62
D. Buskulic,
7
C. Buy,
30
R. L. Byer,
40
L. Cadonati,
63
G. Cagnoli,
64
,
65
C. Cahillane,
1
J. Calder ́on Bustillo,
66
,
63
T. Callister,
1
E. Calloni,
67
,
4
J. B. Camp,
68
K. C. Cannon,
69
J. Cao,
70
C. D. Capano,
8
E. Capocasa,
30
F. Carbognani,
34
S. Caride,
71
J. Casanueva Diaz,
23
C. Casentini,
25
,
13
S. Caudill,
16
M. Cavagli`a,
21
F. Cavalier,
23
R. Cavalieri,
34
G. Cella,
19
C. B. Cepeda,
1
L. Cerboni Baiardi,
57
,
58
G. Cerretani,
18
,
19
E. Cesarini,
25
,
13
R. Chakraborty,
1
T. Chalermsongsak,
1
S. J. Chamberlin,
72
M. Chan,
36
S. Chao,
73
P. Charlton,
74
E. Chassande-Mottin,
30
H. Y. Chen,
75
Y. Chen,
76
C. Cheng,
73
A. Chincarini,
47
A. Chiummo,
34
H. S. Cho,
77
M. Cho,
62
J. H. Chow,
20
N. Christensen,
78
Q. Chu,
51
S. Chua,
60
S. Chung,
51
G. Ciani,
5
F. Clara,
37
J. A. Clark,
63
F. Cleva,
53
E. Coccia,
25
,
12
,
13
P.-F. Cohadon,
60
A. Colla,
79
,
28
C. G. Collette,
80
L. Cominsky,
81
M. Constancio Jr.,
11
A. Conte,
79
,
28
L. Conti,
42
D. Cook,
37
T. R. Corbitt,
2
N. Cornish,
31
A. Corsi,
71
S. Cortese,
34
C. A. Costa,
11
M. W. Coughlin,
78
S. B. Coughlin,
82
J.-P. Coulon,
53
S. T. Countryman,
39
P. Couvares,
1
E. E. Cowan,
63
D. M. Coward,
51
M. J. Cowart,
6
D. C. Coyne,
1
R. Coyne,
71
K. Craig,
36
J. D. E. Creighton,
16
J. Cripe,
2
S. G. Crowder,
83
A. Cumming,
36
L. Cunningham,
36
E. Cuoco,
34
T. Dal Canton,
8
S. L. Danilishin,
36
S. D’Antonio,
13
K. Danzmann,
17
,
8
N. S. Darman,
84
V. Dattilo,
34
I. Dave,
48
H. P. Daveloza,
85
M. Davier,
23
G. S. Davies,
36
E. J. Daw,
86
R. Day,
34
D. DeBra,
40
G. Debreczeni,
38
J. Degallaix,
65
M. De Laurentis,
67
,
4
S. Del ́eglise,
60
W. Del Pozzo,
45
T. Denker,
8
,
17
T. Dent,
8
H. Dereli,
53
V. Dergachev,
1
R. T. DeRosa,
6
R. De Rosa,
67
,
4
R. DeSalvo,
87
S. Dhurandhar,
14
M. C. D ́ıaz,
85
L. Di Fiore,
4
M. Di Giovanni,
79
,
28
A. Di Lieto,
18
,
19
S. Di Pace,
79
,
28
I. Di Palma,
29
,
8
A. Di Virgilio,
19
G. Dojcinoski,
88
V. Dolique,
65
F. Donovan,
10
K. L. Dooley,
21
S. Doravari,
6
,
8
R. Douglas,
36
T. P. Downes,
16
M. Drago,
8
,
89
,
90
R. W. P. Drever,
1
J. C. Driggers,
37
Z. Du,
70
M. Ducrot,
7
S. E. Dwyer,
37
T. B. Edo,
86
M. C. Edwards,
78
A. Effler,
6
H.-B. Eggenstein,
8
P. Ehrens,
1
J. Eichholz,
5
S. S. Eikenberry,
5
W. Engels,
76
R. C. Essick,
10
T. Etzel,
1
M. Evans,
10
T. M. Evans,
6
R. Everett,
72
M. Factourovich,
39
V. Fafone,
25
,
13
,
12
H. Fair,
35
S. Fairhurst,
91
X. Fan,
70
Q. Fang,
51
S. Farinon,
47
B. Farr,
75
W. M. Farr,
45
M. Favata,
88
M. Fays,
91
H. Fehrmann,
8
M. M. Fejer,
40
I. Ferrante,
18
,
19
E. C. Ferreira,
11
F. Ferrini,
34
F. Fidecaro,
18
,
19
I. Fiori,
34
D. Fiorucci,
30
R. P. Fisher,
35
R. Flaminio,
65
,
92
M. Fletcher,
36
J.-D. Fournier,
53
S. Franco,
23
S. Frasca,
79
,
28
F. Frasconi,
19
Z. Frei,
54
A. Freise,
45
R. Frey,
59
V. Frey,
23
T. T. Fricke,
8
P. Fritschel,
10
V. V. Frolov,
6
P. Fulda,
5
M. Fyffe,
6
H. A. G. Gabbard,
21
J. R. Gair,
93
L. Gammaitoni,
32
,
33
S. G. Gaonkar,
14
F. Garufi,
67
,
4
A. Gatto,
30
G. Gaur,
94
,
95
N. Gehrels,
68
G. Gemme,
47
B. Gendre,
53
E. Genin,
34
A. Gennai,
19
J. George,
48
L. Gergely,
96
V. Germain,
7
Archisman Ghosh,
15
S. Ghosh,
52
,
9
J. A. Giaime,
2
,
6
K. D. Giardina,
6
A. Giazotto,
19
K. Gill,
97
A. Glaefke,
36
E. Goetz,
98
R. Goetz,
5
L. Gondan,
54
G. Gonz ́alez,
2
J. M. Gonzalez Castro,
18
,
19
A. Gopakumar,
99
N. A. Gordon,
36
M. L. Gorodetsky,
49
S. E. Gossan,
1
M. Gosselin,
34
R. Gouaty,
7
C. Graef,
36
P. B. Graff,
62
M. Granata,
65
A. Grant,
36
S. Gras,
10
C. Gray,
37
G. Greco,
57
,
58
A. C. Green,
45
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52
H. Grote,
8
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29
G. M. Guidi,
57
,
58
X. Guo,
70
A. Gupta,
14
M. K. Gupta,
95
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1
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1
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