arXiv:gr-qc/0308043v3 17 Sep 2003
Detector Description and Performance
for the First Coincidence Observations
between LIGO and GEO
The LIGO Scientific Collaboration
Corresponding Author: David Shoemaker, MIT NW17-161, 175 Albany St., Cambridge, MA 02139
Tel: 617 253 6411Fax: 617 253 7014Email: email@example.com
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Preprint submitted to Elsevier Science4 February 2008
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J. Loganm,18, M. Lormandp, M. Lubinskio, H. L¨ uckae,b, T. T. Lyonsm,18,
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aAlbert-Einstein-Institut, Max-Planck-Institut f¨ ur Gravitationsphysik, D-14476
bAlbert-Einstein-Institut, Max-Planck-Institut f¨ ur Gravitationsphysik, D-30167
cAustralian National University, Canberra, 0200, Australia
dCalifornia Institute of Technology, Pasadena, CA 91125, USA
eCalifornia State University Dominguez Hills, Carson, CA 90747, USA
fCaltech-CaRT, Pasadena, CA 91125, USA
gCardiff University, Cardiff, CF2 3YB, United Kingdom
hCarleton College, Northfield, MN 55057, USA
iCornell University, Ithaca, NY 14853, USA
jFermi National Accelerator Laboratory, Batavia, IL 60510, USA
kHobart and William Smith Colleges, Geneva, NY 14456, USA
ℓInter-University Centre for Astronomy and Astrophysics, Pune - 411007, India
mLIGO - California Institute of Technology, Pasadena, CA 91125, USA
nLIGO - Massachusetts Institute of Technology, Cambridge, MA 02139, USA
oLIGO Hanford Observatory, Richland, WA 99352, USA
pLIGO Livingston Observatory, Livingston, LA 70754, USA
qLouisiana State University, Baton Rouge, LA 70803, USA
rLouisiana Tech University, Ruston, LA 71272, USA
sLoyola University, New Orleans, LA 70118, USA
tMax Planck Institut f¨ ur Quantenoptik, D-85748, Garching, Germany
uNASA/Goddard Space Flight Center, Greenbelt, MD 20771, USA
vNational Astronomical Observatory of Japan, Tokyo 181-8588, Japan
wNorthwestern University, Evanston, IL 60208, USA
xSalish Kootenai College, Pablo, MT 59855, USA
ySoutheastern Louisiana University, Hammond, LA 70402, USA
zStanford University, Stanford, CA 94305, USA
aaSyracuse University, Syracuse, NY 13244, USA
abThe Pennsylvania State University, University Park, PA 16802, USA
acThe University of Texas at Brownsville and Texas Southmost College,
Brownsville, TX 78520, USA
adTrinity University, San Antonio, TX 78212, USA
aeUniversit¨ at Hannover, D-30167 Hannover, Germany
afUniversitat de les Illes Balears, E-07071 Palma de Mallorca, Spain
agUniversity of Birmingham, Birmingham, B15 2TT, United Kingdom
ahUniversity of Florida, Gainsville, FL 32611, USA
aiUniversity of Glasgow, Glasgow, G12 8QQ, United Kingdom
ajUniversity of Michigan, Ann Arbor, MI 48109, USA
akUniversity of Oregon, Eugene, OR 97403, USA
aℓUniversity of Rochester, Rochester, NY 14627, USA
amUniversity of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA
anWashington State University, Pullman, WA 99164, USA
For 17 days in August and September 2002, the LIGO and GEO interferometer
gravitational wave detectors were operated in coincidence to produce their first
data for scientific analysis. Although the detectors were still far from their design
sensitivity levels, the data can be used to place better upper limits on the flux of
gravitational waves incident on the earth than previous direct measurements. This
paper describes the instruments and the data in some detail, as a companion to
analysis papers based on the first data.
Key words: LIGO, gravitational wave, interferometer, observatory
PACS: 04.89.Nn, 07.60.Ly, 95.45.+i, 95.55.Ym
A number of laboratories around the world [TAMA, VIRGO, GEO,,
LIGO[5,6]] are developing detectors for gravitational waves using laser in-
terferometers to sense the very small strains anticipated from astrophysical
sources. In a joint effort, two of these laboratories, LIGO and GEO600, have
performed their first scientific observations. This note is intended to provide
greater detail in the description of the detectors themselves as a companion
to papers describing the data analysis and astrophysical conclusions from this
Science Run (designated S1).
Both GEO600 and LIGO use the principle of the Michelson interferometer,
1Currently at Ball Aerospace Corporation
2Currently at University of Delaware
3Currently at European Commission, DG Research, Brussels, Belgium
4Currently at European Gravitational Observatory
5Currently at ESA Science and Technology Center
6Currently at Harvard University
7Currently at Hofstra University
8Currently at HP Laboratories
9Currently at Institute of Advanced Physics, Baton Rouge, LA
10Currently at Intel Corp.
11Currently at NASA Jet Propulsion Laboratory
12Currently at Keck Observatory
13Currently at Laboratoire d’Annecy-le-Vieux de Physique des Particules
14Currently at LightBit Corporation
15Currently at Lightconnect Inc.
16Currently at Lockheed-Martin Corporation
17Currently at Laser Zentrum Hannover
18Currently at Mission Research Corporation
19Currently at NASA Goddard Space Flight Center
20Currently at National Science Foundation
21Currently at Rutherford Appleton Laboratory
22Currently at Raytheon Corporation
23Currently at Research Electro-Optics Inc.
24Currently at University of Chicago
25Currently at University of Sheffield
26Currently at Siemens AG
27Currently at Shanghai Astronomical Observatory
28Currently at Stanford Linear Accelerator Center
29Currently at Spectra Physics Corporation
30Currently at University of California, Los Angeles
31Currently at Carl Zeiss GmbH
32Permanent address: University College Dublin
33Permanent address: GReCO, Institut d’Astrophysique de Paris (CNRS)
34Permanent address: University of Tokyo, Institute for Cosmic Ray Research
with its high sensitivity to differential changes ∆L = L1−L2of the lengths L
in the two perpendicular arm lengths L1and L2, to detect strains of the order
of h = ∆L/L = 10−20over a wide frequency range. The required sensitivity of
the interferometric readout is achieved through the use of high circulating laser
power (to improve the shot-noise limited fringe resolution) and through tech-
niques to store the light in the interferometer arms (to increase the phase shift
due to a passing gravitational wave). The frequency range of interest for these
instruments lies in the audio band (∼50-5000 Hz), leading to gravitational
wavelengths λgw= c/fgwof hundreds of km. Because practical ground-based
detectors are short compared to the wavelength, long interferometer arms are
chosen to increase the sensitivity of the instrument. A vacuum system protects
the beams from variations in the light path due to air density fluctuations.
The test masses, which also serve as mirrors for the Michelson interferometer,
are suspended as pendulums and respond as free masses above their ∼1 Hz
resonant frequency. External mechanical disturbances are suppressed through
seismic isolation systems, and the in-band intrinsic thermal noise is controlled
via careful choice of materials and construction techniques.
The LIGO Observatory construction started in 1994 at the LIGO site in Han-
ford, Washington, USA. Construction at the Livingston, Louisiana USA site
began a year later in June 1995. The buildings and 4 km concrete arm-support
foundations were completed in 1998. The vacuum systems were completed in
1999, and detector installation was substantially completed in 2000. The first
operation of a LIGO interferometer took place in October 2000. This marked
the initiation of the commissioning, consisting of periods of intense testing and
tuning of subsystems, separated by periods where the interferometers were run
as complete systems. These engineering runs were primarily intended to assess
the progress toward full detector operation. However, they were also used to
collect data that could be used to test data handling, archiving and analysis
software. Progress through the commissioning phase has been steady, both in
terms of improving sensitivity and in terms of reliable operation with a rea-
sonable duty cycle. By summer 2002, the improvements had been sufficient
that a short duration Science Run could be expected to achieve limits on the
observations of gravitational waves that would be comparable to or better
than previous experimental limits. Consequently, a two-week observation pe-
riod was scheduled, and other laboratories operating interferometer detectors
were invited to join in simultaneous operation, as documented here. Further
progress in sensitivity has subsequently been achieved through additional com-
The construction of GEO600 started in 1995 as a German/British collabo-
ration on a site near Hannover, Germany. Because the site constrained the
length of the arms to 600m, an advanced optical layout and novel techniques
for the suspension systems were included in the detector design. After the
buildings and the trenches were finished in 1997 the complete vacuum system
was installed and tested. The construction phase was followed by the installa-
tion of the two suspended triangular mode cleaners which have been operating
reliably since 2000. To gain experience with the alignment and length control
of long baseline cavities the commissioning continued with the installation
of a 1200 m long Fabry-Perot cavity formed by one interferometer arm and
the power recycling mirrors. To reduce the risk of contaminating or damaging
the expensive main interferometer mirrors, lower grade test mirrors suspended
in steel wire slings were used for the 1200 m cavity experiment and for the
commissioning of the power recycled Michelson interferometer which started
in summer 2001. A first engineering test run was conducted in Jan 2002 in
coincidence with a LIGO engineering run.
The installation of the automatic alignment system for the Michelson interfer-
ometer and for the power-recycling cavity was a key step towards a duty cycle
of more than 98% which was achieved in the 17 day S1 Science Run. The in-
strument ran as a power-recycled folded-arm Michelson for S1; commissioning
of signal recycling started after S1 and is expected to bring the GEO detector
a significant step closer to its design sensitivity.
2Purpose of the S1 Science Run
The primary goal of the S1 run was to collect a significant body of data to
analyze for gravitational waves. Although the sensitivity of all the instruments
was far from the design goal and the relatively short run time made it unlikely
that a positive detection would be made, it was expected that upper limits
could be derived from the data that would be comparable to or better than
previous gravitational wave observations. Furthermore, the analysis provided
the opportunity to test the methodologies with real gravitational wave detec-
tor data. Estimates of sensitivity for gravitational wave interferometers have
almost always been based on the assumption of Gaussian noise. While this is
a good point of departure for many of the limiting noise sources (e.g., shot
noise or thermal noise), many others (e.g., seismic noise) are not expected to
be so well-behaved. Thus, letting the data analysis confront the behavior of
real noise as early as possible is crucial to developing and testing the analysis
Other goals for S1 were aimed at improving our understanding of the detectors
and their operation. These include:
(1) Investigating the factors that influence duty cycle for the interferometers.
Long periods of operation with stable conditions are important for under-
standing the causes for the interferometers to lose ‘lock’ (loss of resonance
condition for light in the interferometer cavities and consequent loss of
linear operation of the sensing system)
(2) Characterization of drifts in alignment and optimization of the alignment
(3) Testing and optimization of on-line monitoring tools for assessing perfor-
mance and maintaining high sensitivity
(4) Training and practice for instrument operators and scientific monitors.
This paper provides a description of the LIGO and GEO interferometers as
they were used in the S1 run.1It is intended as a companion to the data
analysis papers based on data from this run. Because commissioning was still
underway, many parts of the detectors were not in their intended final op-
erational configuration, and an important emphasis of this paper will be to
identify and highlight those differences.
3 The LIGO detector array
The LIGO detector array comprises three interferometers at two sites. The
LIGO Livingston Observatory (LLO) contains three main instrument bays at
the vertex and ends of the L-shaped site and houses a single interferometer
with 4 km long arms (designated L1). The LIGO Hanford Observatory (LHO)
has two additional experimental halls at the midpoint in each arm which
enable it to accommodate two interferometers, one with 4 km long arms (des-
ignated H1) and one with 2 km arms (H2). The orientation of the Hanford
site was chosen to be as closely aligned (modulo 90◦) to the Livingston site as
possible, consistent with the earth’s curvature and the need for the sites to be
level; this maximizes the common response to a signal, given the quadrupolar
form of the anticipated gravitational waves. The arms have an included angle
of 90.000◦. The locations and orientations of the two LIGO sites are given in
The observatories have a support infrastructure of clean rooms, preparation
laboratories, maintenance shops, and computer networking for control, data
1A shorter period of simultaneous observations between TAMA, GEO, and LIGO
also took place during the period of this science run. That effort will be documented
acquisition and analysis, as well as offices for site staff and meeting spaces for
larger gatherings. The vacuum system can be divided into two main pieces: the
4 km beam tube arms (through which the laser beams pass between the ver-
tex and end test masses), and the vacuum chambers that house the suspended
optics and associated equipment. The vacuum tubing for the arms[13,14] is
1.2 m in diameter, fabricated from 3 mm thick 304L stainless steel, processed
to reduce the outgassing to very low levels (1 to 8×10−14mbar · L · s−1cm−2).
Expansion bellows are placed periodically along the arms. An extended bake
at elevated temperature was used to remove adsorbed water. The tubing is
supported by a ground-level concrete slab, protected by a concrete cover, and
aligned to centimeter accuracy. Ion pumps at 2 km intervals and liquid
nitrogen cooled cryogenic traps where the arms enter the buildings at the ver-
tex, end, and midstations maintain the base pressure in the arms between 10−8
and 10−9mbar, with the residual gas being mainly hydrogen. This pressure
is sufficient to put the residual gas scintillation well below the LIGO design
The seismic isolation system, test masses, and other interferometer optics are
housed in vacuum chambers at the vertex, mid-stations (at Hanford), and
end-stations. Large gate valves where the beam tubes enter the buildings al-
low the vacuum chambers to be isolated from the beam tubes and brought
to atmospheric pressure for work on the suspended optics while maintaining
the vacuum in the 4 km arms. The vacuum chambers have large doors to
aid in the access to install and align optics. When the chambers are at at-
mospheric pressure they are purged continuously with clean (Class 10) dry
air. They have numerous viewports (for laser beams to enter and exit the
vacuum system and for video camera monitoring of the interior components)
and electrical feedthroughs. The pumping system includes roughing pumps
with Roots blowers, and hydrocarbon-free turbopumps. Only ion pumps and
cryogenic traps are used when the interferometers are operating. The vacuum
chambers are fully instrumented with gauges and residual gas analyzers; pres-
sures range between 4 × 10−8and 3 × 10−9mbar. All materials used in the
vacuum chambers and for the installed detector equipment are carefully pro-
cessed and screened to minimize the amount of hydrocarbons introduced into
Location and orientation of the LIGO detectors. Note that the Livingston Observa-
tory is rotated by ∼ 90◦with respect to the Hanford Observatory, when the earth’s
curvature is taken into account.
Vertex Latitude46◦27’18.5” N30◦33’ 46.4” N
Vertex Longitude119◦24’ 27.6” W90◦46’ 27.3” W
Orientation of X arm324.0◦(NW)252.3◦(WSW)
Fig. 1. Schematic layout of a LIGO interferometer.
the vacuum system as a precaution against mirror contamination.
The basic optical configuration of each LIGO detector is that of a power-
recycled Michelson interferometer with resonant arm cavities, shown in Fig-
ure 1. Gravitational waves produce strains in space. The arm cavity mirrors
serve as the inertial test bodies (test masses), which move in response to these
strains. For example a sinusoidal wave incident on the plane of the interfer-
ometer will cause an apparent shortening of the optical path along one arm
and a lengthening along the other arm, and this process reverses half a cycle
later in the signal evolution. Laser light is incident from the bottom-left on
the beamsplitter, which divides it and sends it to low-loss cavities in the arms.
The transmission of the input mirror in each cavity is much larger than the
losses in the cavity, and thus when the cavities are on resonance, almost all
of the light is returned to the beamsplitter. The beamsplitter is held so that
the light emerging from the antisymmetric port of the interferometer (right)
is at a minimum, and almost all of the light is reflected back toward the laser.
The power-recycling mirror forms a resonant optical cavity with the interfer-
ometer, causing a build-up of power in the recycling cavity. The arm cavity
mirrors serve as the inertial test bodies (test masses), moving in response to
the gravitational wave.
Each interferometer is illuminated with a continuous-wave Nd:YAG laser op-
erating in the TEM00Gaussian spatial mode at 1064 nm, and capable of 10
W output power. A small portion of the beam is used to stabilize the
laser frequency using a reference cavity in an auxiliary vacuum chamber (Fig-
ure 2). The beam going to the reference cavity is double-passed through an
acousto-optic modulator driven by a voltage controlled oscillator; this allows
an offset between the laser frequency and the reference cavity frequency to
permit the laser to follow the arm cavity length change due to tidal strains.
This initial level of stabilization is at 0.1 Hz/Hz1/2or better in the gravita-
tional wave band. The main portion of the beam is passed through a 45 cm
path length triangular cavity to strip off non-TEM00light and to provide pas-
sive filtering of the laser intensity noise with a pole frequency of 1.5 MHz (0.5
MHz at Livingston for the S1 run). An intensity noise control system designed
to reduce relative intensity fluctuations below 10−8Hz−1/2was only partially
implemented during the S1 run, leaving the intensity noise at approximately
10−7Hz−1/2level. Electro-optic modulators impress radio-frequency sidebands
on the light at 24.5 and 33 MHz (29.5 and 26.7 MHz for H2) for sensing respec-
tively the interferometer, and suspended mode cleaner degrees of freedom.
The design for the LIGO interferometers is for 8 W to be incident on the
mode cleaner. However, the commissioning of the instrument for high input
power was not completed at the time of S1, and the powers incident on the
mode cleaner had been adjusted (through the use of attenuators and reduced
laser power) to approximately 1 W for H1 and L1 and approximately 6 W for
3.2 Input Optics
After the laser beam enters the vacuum system, it passes through a set of
input optics to condition it before it passes to the main interferometer. First,
it passes through a mode cleaner – a ∼24 m path length triangular ring cavity
with a finesse of ∼1350, formed from separately suspended mirrors. This cavity
stabilizes the beam size, position and pointing. It also blocks the 33 (or 26.7)
MHz sidebands, but transmits the 24.5 (or 29.5) MHz sidebands used for the
interferometer sensing, which are at multiples of the mode cleaner free spectral
range. In addition, it serves as an auxiliary reference for the laser frequency
control servo, reducing frequency noise in the transmitted laser light to the
10−3Hz/Hz1/2level for the S1 run parameters. After the mode cleaner, the
beam passes through a Faraday isolator, which diverts light returning from the
interferometer onto a photodetector. This prevents the returning light from
reaching the laser and causing excess noise, and allows the common-mode
motions of the test-mass mirrors to be sensed. Finally, the beam passes through
an off-axis telescope formed by three suspended mirrors, which expands the
beam to match the ∼4 cm (1/e2radius in power) mode of the arm cavities.
Fig. 2. Simplified schematic of laser stabilization. EOM: Electro-Optic Modulator;
AOM: Acousto-Optic Modulator; VCO: Voltage Controlled Oscillator; PD: Photo
Diode; PMC: Pre-Mode Cleaner; IOO: Input Optics; LSC: Length Sensing/Control
input (from IOO)
input (from LSC)
to input optics
pwr. stab. input
pwr. stab. amp.
tidal stab. amp.
freq. stab. amp.
3.3 Interferometer Optics
The main interferometer optics[19,20] are fabricated from high-purity fused
silica, 25 cm in diameter and 10 cm thick (except the beamsplitter which is
4 cm thick). Radii of curvature of the cavity optics are chosen so that the
arm cavities have a stability g = (1 − L/R1)(1 − L/R2) (L is the cavity
length and Rnare the radii of curvature of the two cavity mirrors) of 0.33 (H1
and L1) or 0.67 (H2), to minimize the excitation of higher order transverse
modes by separating them in frequency from the laser frequency and its RF
modulation sidebands. The surface figure accuracy of the polished optic is
better than 1 nm; the coatings have a thickness uniformity that holds their
contribution to the apparent surface flatness negligible. The coatings have a
power absorption less than 1 ppm and scatter less than 70 ppm. All optics are
wedged (typically about 2 degrees) to reduce the possibility of stray reflections
interfering with the main beam and to give access to samples of the light inside
the interferometer. Transmission of the input mirrors to the arm cavities is
2.7% and the end mirrors have a transmission of approximately 12 ppm, to
give an arm cavity pole frequency of 85 Hz ( 170 Hz for H2). The beamsplitter
reflectivity was specified as 50±0.5%. The recycling mirror transmission is also
2.7%, to give a design recycling factor (or increase in the circulating power)
of ∼50 for the optics as designed and at full power.
During the S1 run, the low light input power led to the optical configurations
of the three interferometers operating away from their design point. At full
operating power, absorption of light in the substrate and coating of the input
mirrors for the arm cavities is expected to create significant thermal lensing.
As a result, the curvature of the recycling mirrors was figured to compensate
for this anticipated thermal lensing. Since the incident laser power in the H1
and L1 interferometers was significantly under the design level, the lack of
thermal lensing makes the recycling cavities slightly unstable for the modula-
tion sideband light. This has little effect on the carrier recycling gain (since the
carrier spatial mode is stabilized by resonance in the arm cavities) but reduces
the transmission of sideband light to the antisymmetric port by more than a
factor of 10. This further reduces the main differential arm length sensitivity
in the high frequency region where shot noise is expected to be dominant.
In the case of the 2 km interferometer H2, although it was receiving nearly
the design input laser power, an out-of-specification anti-reflection coating on
the input mirror of one arm caused excess loss in the recycling cavity and
reduced the recycling gain for the carrier by more than a factor of two. As a
result it also did not develop the required thermal lens and its transmission
of sidebands to the dark port was also degraded by a similar factor. These
limitations contributed to the relatively high noise level of the instruments
seen in the sensing-noise limited regime (f > 200 Hz).
Each interferometer optic is suspended as a pendulum from vibration-isolated
platforms to attenuate external disturbances in the gravitational wave band;
see Figure 3 for a schematic drawing. The suspension fiber is a steel piano
wire, loaded at approximately 40% of its yield stress, passing under the optic
as a simple loop. Small, notched glass rods are glued to the side of the optic
a few millimeters above the center of mass to define the suspension point and
minimize frictional losses. The normal modes of the test mass optic suspen-
sion are approximately 0.74 Hz (pendulum mode), 0.5 Hz (yaw mode), 0.6
Hz (pitch mode), 12 Hz (bounce mode), 18 Hz (roll mode) and multiples of
345 Hz (violin mode). Thermal noise is managed in interferometric gravita-
tional wave detectors by placing resonances above or below the detection band
when possible, and by choosing materials and assembly techniques which yield
high resonance Q’s. This gathers the thermal noise power into a narrow
band and lowers the values on either side of the resonance. In the case of the
suspensions, high resonance Q’s (measured to be typically 2 to 4×105) in all
suspension modes yield a negligible level of off-resonance thermal noise for the
The suspension system also provides the means for applying control forces
and torques to align the mirrors and hold the interferometer in resonance.
Four small Nd:Fe:B magnets are attached to the back of the mirror using
aluminum stand-offs and a vacuum compatible epoxy, with alternating po-
larities to reduce coupling to environmental magnetic fields. The suspension
structure supports voice coils on ceramic forms near the magnets to produce
control forces. Each of these assemblies also incorporates an LED/photodiode
pair arranged so that the magnet partially obstructs the path between them
Fig. 3. LIGO Suspension
(a “shadow sensor”). This provides a read-out of the longitudinal position of
the magnet with a noise level of approximately 10−10m/Hz1/2in the gravita-
tional wave band. Similar magnets are attached to the sides of the optic and
a shadow sensor/voice coil assembly acts on one of these to damp sideways
The magnet and coil actuators are driven by several sensors, via servo con-
trollers, allowing control of their positions and orientations with respect to
both the local structures and the globally-measured lengths and angles. Local
damping of the modes of the suspension is provided by feeding appropriately
filtered and mixed signals from the shadow sensors to the coils to create a
damping force near the pendulum frequencies. Signals from interferometric-
based wavefront sensors and optical levers (described below) are also applied
to maintain the pointing of the test mass. Lastly, the interferometer length sig-
nals are applied to acquire resonance and hold the operating lengths for the
interferometer to within ∼ 10−13m rms. The suspension controllers, which
combine and filter these signals appropriately, were of two styles during S1:
an original analog system with some digital gain and filter controls (for H2 and
L1), and a system with the signal processing performed digitally (H1). In
all cases, a significant low-frequency noise contributor was the final amplifier,
which for S1 had to deliver stronger control forces than those expected for
the final configuration. This then compromised the gravitational wave-band
Thermal noise internal to the mirrors is minimized by maintaining high Q’s in
all the internal modes. The fused silica internal losses are anticipated to make
the dominant contribution to thermal noise. However, the dielectric coating
on the mirror will also contribute noticeably because of its proximity to the
beam. The attachments to the mirror for the suspension and the magnets
can degrade the individual modal Q’s but, because of their distance from the
front surface of the optic, their effect on thermal noise is negligible. In-situ
measurements of Q’s typically range from 2 × 105to 1.6 × 107, depending on
the mode. Calculations indicate that the thermal noise is near the design goal
and thus negligible for the S1-run sensitivity.
3.5 Seismic Isolation
The vibration isolation systems are four layer passive isolation stacks. The
final stage in each vacuum chamber is an aluminum optical table that holds
the optic suspensions. Each optical table is supported by four legs. Each leg
consists of a series of three heavy stainless steel cylinders, supported by coil
springs made with phosphor bronze tubing containing inner constrained layers
which are sealed from the vacuum via electron-beam welding. The transfer
function of ground motion to table motion shows a series of broad peaks
between 1.5 and 12 Hz, representing the normal modes (typical Q ∼ 10 −
30) of the masses moving on the springs, followed by a steep falloff above
the highest resonance. The total attenuation reaches a value of 106by about
50 Hz. The high Q’s of the resonances in the 1.5 to 12 Hz band presents a
particular problem at LLO, where they amplify anthropogenic ground noise
in this frequency range, and cause difficulties in locking during daylight hours.
A planned six degree-of-freedom external active isolation system to cope with
this excitation was not in place during S1. The support points for the seismic
isolation stack penetrate the vacuum chamber through bellows that decouple
the seismic isolation stack from vacuum chamber vibrations and drift. External
coarse actuators at the support points permit translations and rotations to
minimize the control forces that are needed to align the optics during the
initial installation and to compensate for any long-term settling.
In addition, the systems at the ends of the arms are equipped with a fine
actuator aligned with the arm that can translate the entire assembly (seis-
mic isolation stack and optic suspension) by approximately ±90µm over the
frequency range from DC to ∼ 10 Hz. This system is used during the inter-
ferometer operation to compensate for earth tides, using a simple predictive
model and a very slow feedback from the differential and common mode arm
length controls. At LLO, an additional microseismic feed-forward system
was used to reduce the length fluctuations of the arms at the microseismic fre-
quency (approximately 0.16 Hz). Also, the L1 detector’s fine actuators were
used together with seismometers in a beam-direction active seismic isolation
system at each test mass chamber, which reduced seismic excitation of the
most troublesome stack modes by a factor of ∼5.
3.6 Length and Angle Control
There are four longitudinal degrees of freedom that must be held to allow the
interferometer to function: the two arm lengths are held at the Fabry-Perot
cavity resonance condition, the beamsplitter position is set to maintain the
light intensity minimum at the antisymmetric port and the recycling mirror
position is positioned to meet the resonance condition in the recycling cavity.
These lengths are sensed using RF phase modulation sidebands on the incident
light in an extension of the Pound-Drever-Hall technique. The modulation
frequency was chosen so that the phase modulation sidebands are nearly an-
tiresonant in the arm cavities; the carrier light is strongly overcoupled so that
0.97 of the light is reflected on resonance, and it receives a π phase shift on
reflection. By making the recycling round trip cavity length an odd number
of RF half-wavelengths, the recycling cavity can be simultaneously resonant
for the carrier and sidebands. A small length asymmetry (∼30 cm) is intro-
duced between the beamsplitter and the two input test masses to couple the
sideband light out the dark port.
Three interferometer output beams (Figure 4) are used to determine the lon-
gitudinal degrees of freedom, which are best thought of as two differen-
tial motions (arm cavities or strain readout, and the Michelson), and two
common-mode motions (common mode ‘breathing’ of the arm cavities, and of
the power recycling cavity). A photodiode signal at the antisymmetric port is
demodulated with the 90◦quadrature of the modulation drive to give a signal
proportional to the difference in arm cavity lengths (differential arm length).
A second photodiode monitors the light reflected from the recycling mirror
(separated from the incident beam by a Faraday isolator); it is demodulated
in phase with the modulation drive and is primarily sensitive to the average
of length of the two arm cavities (common mode arm length). The third pho-
todiode monitors light from inside the recycling cavity, picked off from the
back (anti-reflection-coated) side of the beamsplitter with the aid of the small
wedge angle in the substrate. The in-phase signal is primarily sensitive to the
recycling cavity length, while the quadrature phase is sensitive to the Michel-
son path difference from the beamsplitter to the input test masses. One major
deviation from the final interferometer design during S1 was that attenuators
were placed in front of the antisymmetric port photodiodes on all three inter-
ferometers, reducing the effective power used in each interferometer to about
50 mW instead of the 6 W nominal value. These attenuators protected the
photodiodes from saturation, and possible damage, during the commission-
ing phase before the complete mirror angular controls were implemented and
when large fluctuations in the power on the photodiodes were present. This
had a particularly significant impact on the performance in the high frequency
region (above a few hundred Hz), where the low effective light level combined
with the reduced sideband efficiency noted above to cause the noise to be well
Fig. 4. Schematic drawing of a LIGO interferometer showing laser, input light mode
cleaner, and the locations of the photodiodes (Sxxx) used to sense and control the
resonance conditions. L1and L2are the arm cavity lengths; a gravitational wave
produces a differential signal of the form (L1− L2), and (L1+ L2) is a sensitive
measure of the laser frequency noise. The Michelson degrees of freedom are differ-
ential (l1− l2) and common-mode (l1+ l2), the latter measured with respect to
the recycling mirror. PBS: Polarizing Beam Splitter. AOM: Acousto-Optic Mod-
ulator. PC: Pockels Cell. VCO: Voltage Controlled Oscillator. Smc: Signal, Mode
Cleaner. Sref: Signal, Reference Cavity. Srefl: Signal, reflected light. FI: Faraday Iso-
lator. PRM: Power Recycling Mirror. Sprc: Signal, Power Recycling Cavity. Santi:
Signal, Antisymmetric Port. BS: Beamsplitter. ITM: Input Test Mass. ETM: End
above the design level.
The signals from these three photodiodes, appropriately demodulated and fil-
tered, are used to control the lengths and hold the interferometer in resonance.
The high frequency portion of the reflected photodiode signal Sreflis fed back
(via an analog path at Hanford for S1, digital at Livingston) to the mode
cleaner and laser to stabilize the input laser frequency to the average arm
length. The signals Sprc and Santi from the other two photodiodes are used
to control the positions of the interferometer optics. The demodulated signals
from all three photodiodes are whitened with an analog filter, digitized with
a 16 bit ADC operating at 16384 samples per second, and digitally filtered
with the inverse of the analog whitening filter to return them to their full
dynamic range. A dedicated real-time signal processor combines these error
signals via a matrix (whose coefficients are adjusted in real time during the
lock acquisition process) to form appropriate control signals, filters them, and
sends the results to combinations of optics to control the interferometer. It
also passes the photodiode signals (error signals) and the feedback signals to
the data acquisition system. The flexibility of the digital control system to re-
spond in changes to the interferometer response function during the ‘locking’
process as a function of sensed light levels, and to allow specialized filters to be
implemented on the fly, has been crucial to the ability to acquire lock on the
interferometers, to aid in the commissioning, and ultimately to suppress
noise in the control systems.
As noted above, an ensemble of optical levers and wavefront sensors is de-
signed to sense and control the angular degrees of freedom of the suspended
optics in the main part of the interferometer[29,30]. Each large (25 cm) optic
is equipped with an optical lever, consisting of a fiber-coupled diode laser and
a (position sensitive) quadrant photodiode, which is intended to hold the optic
stable while the interferometer is unlocked. These components are mounted on
piers outside the vacuum system and operate through viewports at distances
between 1 and 25 meters from the optic; their long-term stability and inde-
pendence from the interferometric sensing system allows a manual alignment
to be maintained continuously.
The full instrument design includes a wavefront sensing control system to
optimize the alignment during operation. Quadrant photodiodes are placed at
the output ports of the interferometer, in the near field and (via telescopes)
in the far field. The photocurrents are demodulated as for the length control
system, and sums and differences can be formed to develop a complete set of
alignment information which is then used to control the mirror angles, using
the suspension actuators. However, at the time of the S1 run, this system
was only partially commissioned, and only the mode cleaner and two degrees
of freedom of the interferometer, the differential pitch and yaw of the end
test masses (cavity end mirrors), were controlled by wavefront sensors. As
an interim measure, the incomplete wavefront sensing was complemented by
signals from the optical levers during operation. However, the optical lever
angular sensing noise is much greater than that for the wavefront sensors.
Even after careful control-law shaping, the optical levers remained one of the
principal contributors to the low-frequency noise of the instrument for S1.
Baffles to capture stray light are placed along the 4 km beam tubes, and at
specific places near the optics inside the vacuum chambers, to reduce the pos-
sibility of a scattered beam or one from an intentional wedge in the optics from
recombining with the main beams. Some of the baffles for the final installation
were not in place for the S1 run, but calculations indicate that this should not
have been a source of noise at our present level of sensitivity.
A noise model of the instruments summarizes the limits to the performance of
the interferometer at the time of this science run. The model for the Livingston
detector is shown in Figure 5. In general, the contributions are evaluated by
Fig. 5. A frequency-domain model of the noise sources at the time of the S1 run for
the Livingston (L1) detector. The noise sources, discussed in the text, are assumed to
add in quadrature. The actual noise curve is also shown, along with the performance
expected for the instrument when working at the design level
S1 L1 Noise spectrum, 08/29/02
Optical Lever Servo
Suspension Coil Driver
Shot and Electronic Noise
measuring a source term (e.g., laser frequency noise), and measuring a coupling
function (e.g., the transfer function from an intentional frequency modulation
to the response in the strain channel), and then multiplying these two together
to make a prediction. In some cases, analytical models are used (for example
the mechanical Q of the suspension systems is measured and then used in a
model of the thermal noise contribution). For this model, all the terms are
considered to be independent, and the noises are added as the square root
of the sum of the squares. Many sources of noise have been modelled; this
figure only shows those that limit the present performance. The model explains
the overall instrument noise performance well, and subsequent commissioning
efforts have shown that reductions in the leading noise terms also leads to the
anticipated reduction in the overall instrument noise.
In addition to subsystem dynamics and control models, two simulations played
a significant role in the design and commissioning of the LIGO detectors. The
first is an FFT-based optical propagation code that models the power-
recycled Michelson interferometer with Fabry Perot arms. This code was used
to develop the specifications for the interferometer optics, and has been used
in comparisons with commissioning data to evaluate the performance of the
optics as installed. The second simulation is an end-to-end time-domain sim-
ulation of the LIGO interferometers. This model includes a modal-based
optical propagation, accurate modeling of the electronic feedback, simplified
models for the suspension systems, and typical noise inputs. This model proved
to be invaluable in developing the lock acquisition software.
3.8 Environmental Monitoring
A system of auxiliary sensors is installed at each LIGO site to monitor possible
environmental disturbances. The Physics Environment Monitor system (PEM)
contains seismometers and tiltmeters to monitor low frequency ground distur-
bances, accelerometers and microphones to monitor higher frequency mechan-
ical disturbances, magnetometers to monitor magnetic fields that might affect
the test masses, and monitors of the line power. Sensors are present in all
buildings and near all key sensitive components. They have been used to e.g.,
help identify sources of acoustic and electromagnetic coupling, and to help de-
sign improvements to the apparatus; as the instrument sensitivity improves,
they will be used as veto signals in the astrophysical analyses. Planned cosmic
ray detectors and rf monitors were not operational at the time of the S1 run.
3.9 Control and Data Systems
Supervisory control of the interferometers is accomplished using EPICS (Ex-
perimental Physics and Industrial Control System). EPICS establishes a
communications protocol within a non-hierarchical computer network and pro-
vides an operator interface from networked workstations located in the control
room. Processors distributed in all electronics racks can modify amplifier gains,
offsets, filtering, on/off controls, etc., allowing either manual or automated
(scripted) control of the state of the electronics. The EPICS processors also
interface to analog-to-digital converters to provide monitors for the electron-
ics inputs and outputs. Each interferometer has approximately 5000 EPICS
variables (either control or monitor points). EPICS also provides tools for cap-
turing and restoring the state of the instrument to ensure that this complex
instrument can be reliably brought to a known configuration.
The data acquisition and control system collects signals from the interferom-
eter and from the environment, and delivers signals for the length and angle
controls. VME-based converters and processors are used, and acquisition sys-
tems are placed in the vertex building and the mid- and end-stations. Analog
signals are digitized with 16-bit resolution. Fiber optics are used to link the
instrument racks together, and a shared memory approach allows data to be
collected and shared by a number of systems over the multi-km distances. The
data are collected with 16-bit resolution. The complete data are formatted into
the standard data ‘Frames’ (a format used by all of the interferometric grav-
itational wave community) and initially are stored on spinning media for a
quick ‘look-back’ buffer of roughly two weeks. All data are archived to tapes
for later analysis. Reduced data sets also in the standard Frame format, con-
figured for a given science run, are produced as well; these serve most analysis
It is important that the data acquisition system accurately time-stamp the
data it records. The fundamental timing for both sites is derived from GPS
receivers located at each building (vertex, mid and end). A 222Hz (approxi-
mately 4.2 MHz) clock signal is generated from the GPS as well as a 1 pulse-
per-second synchronization signal. Together these are used to synchronize the
data collected by the various processors. Ramp signals are used to monitor any
timing errors and alarms are set for the operators. This monitoring has proven
useful during S1. It showed that the timing was subject to jumps (typically
10’s of milliseconds, but sometimes larger) when the length control system
processors were rebooted, with the consequence that some S1 data had to
be eliminated from some analyses because of uncertain timing (These timing
jumps have since been cured, and a redundant and independent atomic clock
reference is being implemented for the future).
3.10 Diagnostics and Monitoring
Two closely related systems, the Global Diagnostic System and the Data Mon-
itoring Tool, provide the instrument operators and scientific monitors with
tools for evaluating interferometer performance both during commissioning
and scientific running. The Global Diagnostic System (or GDS) can ac-
cess data from any signal collected by the data acquisition system, including
test-point signals that can be stored for post-analysis if indicated. It can dis-
play the time series and the power spectrum for individual signals, and the
transfer function and coherence for pairs of signals. The GDS also has the abil-
ity to apply arbitrary-waveform excitations to various test points within the
interferometer control systems. These can be used to measure transfer func-
tions through stimulus-response testing. The data can be filtered, decimated,
calibrated, stored and recalled for comparisons.
The Data Monitoring Tool (or DMT) is a package of software components run-
ning on a set of processors on a dedicated network. A high-speed connection
to the data acquisition system makes the full data set available with only a
one-second latency. The emphasis in the DMT is on relatively simple measures
of instrument performance applied to the full data stream in realtime. Thus it
can give the operators and scientific monitors rapid feedback about interfer-
ometer performance. These include such measures as the non-stationarity and
burst-like behavior of various types in the interferometer outputs, band-limited
rms amplitudes of interferometer outputs and environmental monitors, mon-
itors of calibration lines, histograms to monitor the gaussianity of the data,
and real-time estimates of detector sensitivity to neutron star binary inspiral
events. The DMT is also an important element of the data analysis process,
analyzing the auxiliary channels for veto signals in parallel with the strain
3.11 Data Analysis System
To analyze the large volume of data generated, LIGO has developed the LIGO
Data Analysis System (LDAS). The LDAS provides a distributed software en-
vironment with scalable hardware configurations to provide the computational
needs for both on and off site data analysis. The architectural design of the
system is based on the concept of multiple concurrent data analysis pipelines
in which data is fed into the pipeline as it is collected and proceeds down the
pipeline where necessary signal analysis procedures are applied depending on
the particular type of analysis that is being carried out.
LDAS is complemented by the LIGO/LSC Algorithm Library, which is a set
of C-language routines that can run under LDAS or be used independently.
They are carefully vetted to ensure that the algorithms and results are correct.
The LDAS distributed software environment is composed of roughly 12 mod-
ules called LDAS Application Programming Interfaces (APIs), each of which
is a separate process under the Unix operating system. Each module is de-
signed to carry out the multitude of steps associated with each unique pipeline.
For example, one module has computational elements for reading and writing
LIGO channel data in the Frame format, another module has computational
elements for signal processing in the time or frequency domain, and another
module has computational elements necessary to perform parallel analysis
across a cluster of tightly networked CPUs. Upon completion of the data
analysis pipeline, data products and results are stored to disk or inserted into
the LIGO relational database. The database has tables designed to capture
results associated with detector characterization, on line and off line astro-
physical searches and multi-detector analyses. The software can be scaled to
run on a wide variety of computing hardware.
During the first LIGO Science Run, the LDAS at the LIGO Observatories and
data analysis centers located at Caltech and MIT, LDAS at other institutions,
and other configuration-controlled computational systems were operated with
commissioning configurations of the hardware and software. The software was
in the late stage of beta development, having a complete set of modules.
The hardware systems consisted of a complement of servers with tens of ter-
abytes of disk storage for the raw data and the LIGO database, along with
scaled down computation centers with approximately 200 megaflops of aggre-
gate computational performance between them. In its final configuration, the
LDAS hardware will include upgrades to the current servers, and expanded
high performance computation clusters with over two teraflops of aggregate
computational performance. In addition, new tape storage systems will be put
on line which will provide adequate storage at the observatories for six months
of local data and storage for all of the data at Caltech; this is where the data
for the multiple detectors are brought together. The software is expected to
double in performance as we upgrade from beta versions to the first completed
version later this year. In addition, the software is being adapted to support
Grid Computing technology and security protocols allowing for LIGO data
analysis once the computational Grid is deployed in the near future.
3.12 LIGO Data
The full data stream from each of the LIGO interferometers consists of several
thousand channels, recorded at rates from 1 Hz to 16384 Hz with a total data
rate of 5MB/s per interferometer. These channels include EPICS process vari-
ables that define the state of the interferometer, signals from environmental
monitors, signals from auxiliary servos in the interferometer (for example, op-
tical lever signals), as well as the main gravitational wave signal. For the servos
not operational during S1, the corresponding data channels were recorded, but
of course they contain only zeros. The non- gravitational wave data channels
can be used in a number of ways:
(1) They can be used to determine the operational “health” of the interfer-
ometer (how well it was aligned, whether large offsets were present in any
(2) They can be used to regress noise from the main gravitational wave chan-
nel (for example, measurements of the laser frequency noise can be used
to correct the gravitational wave channel to remove any residual effects
from laser frequency noise coupling to mismatches in the arms).
(3) They can be used to veto non-gaussian noise in the interferometer (for
example seismometer data could be used to keep noise from impulsive
seismic disturbances from being misinterpreted as a gravitational wave).
At this point in the commissioning, few of these techniques have been explored
and developed. In part, this is because the majority of the noise sources in the
interferometer at present are attributable to electronic noise entering through
imperfect tuning, and consequently, few of the auxiliary channels are expected
to be useful. The DMT capabilities to perform this analysis were exercised in
preparation and performed very well.
The main signal for the analysis to search for gravitational wave signals is
the output of the photodiode at the antisymmetric port, demodulated in the
quadrature phase at 24.5 MHz (H1 and L1), or 29.5 MHz (H2). This analog
signal is amplified and digitized. An analog filter whitens the signals before
digitization, and a precise inverse of this filter ‘de-whitens’ the signal in the
digital domain, to best take advantage of the dynamic range and noise in
the Analog-to Digital Converter (ADC). Since it is the error point in the
servo control system which holds the differential arm length, its interpretation
requires correction for the loop gain of the servo. This signal represents the
phase difference of the light from the two arms, filtered only by a roll-off
at high frequencies because of the arm cavity storage time, and while the
interferometer is operating, is a continuous measure of the differential strain
between the two arms and thus potentially of gravitational wave signals.
4 The GEO Detector
The GEO Detector is situated at the perimeter of an agricultural research sta-
tion to the south-east of Hannover, Germany; see Table 2. The buildings are
intended to be just sufficient to accommodate the instrument and its acqui-
sition and control hardware. Data recording, and much of the operation and
on-line monitoring of the instrument, will be performed at the Max Planck In-
stitute in downtown Hannover, once continuous science operation is underway.
A microwave link maintains a high-bandwidth dedicated connection between
Location and orientation of the GEO600 detector. Note that the arms form an
angle of 94◦19’ 53”. This deviation from perpendicular has negligible effect on the
Vertex Latitude52◦14’ 42.5” N
Vertex Longitude9◦48’ 25.9” E
Orientation of North arm
Orientation of East arm
One central building (13 m × 8 m in size) and two end buildings (6 m × 3
m) accommodate the vacuum chambers (2 m tall, 1m in diameter) in which
the optical components are suspended. In the central building, nine vacuum
chambers form a cluster which can be subdivided into three sections to allow
mirror installation without venting the whole cluster. This arrangement allows
a minimum of down-time for a change of the signal-recycling mirror (to change
the detector bandwidth). To avoid fluctuations of the optical path caused by
a time-varying index of refraction, all light paths in the interferometer are in a
high-vacuum system. For this purpose GEO600 uses two 600 m long vacuum
tubes of 60 cm diameter which are suspended in a trench under ground. A
novel convoluted-tube design, allowing a wall thickness of only 0.8 mm, was
used to reduce weight and cost of the stainless-steel vacuum tube.
The whole vacuum system, except for the mode cleaner and signal recycling
section, is pumped by four magnetically levitated turbo pumps with a pumping
speed of 1000 l/s, each backed by a scroll pump (25 m3/h). Due to the use of
stainless steel with a low outgassing rate, a 2 day air bake at 200◦C and a 5
day vacuum bake at 250◦C, a pressure of 1 × 10−8mbar can be achieved in
the tubes. Large gate valves allow the beam tubes to be temporarily closed off
and maintained under vacuum whenever the instrument vacuum chambers are
opened for installation work. Additional dedicated pumping systems are used
for the mode cleaner section and for the signal-recycling section. The pressure
in the vacuum chambers is in the mid 10−8mbar range. Great care was taken to
minimize contamination of the all-metal vacuum system by hydrocarbons. For
this reason the seismic isolation stacks, which contain silicone elastomer and
other materials containing hydrocarbons, are sealed by bellows and pumped
separately. Furthermore, the light emitting diodes (LEDs), the photodiodes
and the feedback coils used as ‘shadow’ sensors and actuators in the pendulum
collocated damping and actuation systems are encapsulated in glass.
The buildings of GEO600 are split into three regions with different cleanroom
classes: the so-called gallery where people can visit and staff can work with
normal clothes, the inner section which has a cleanroom class of 1000 and a
movable cleanroom tent installed over open chambers with a cleanroom class
4.1Suspension and Seismic Isolation
Two different types of seismic isolation and suspension systems are imple-
mented in GEO600. The first one, used to isolate the mode cleaner optics,
consists of a double pendulum suspended from a pre-isolated top-plate. To
avoid an excitation of the pendulum mode, four collocated control systems
measure the motion of the intermediate mass with respect to a coil-holder
arm which is rigidly attached to the top plate, and feed back to the mirror via
a coil-and-magnet system.
The seismic isolation system used to isolate the test masses, beamsplitter,
and the other mirrors of the Michelson interferometer consists of a triple
pendulum suspended from a pre-isolated platform. Each pendulum chain
consists of the optic suspended from an intermediate mass which is in turn
suspended from an upper mass. Two cantilever spring stages are included in
the pendulum design (in the support of the upper and intermediate masses) to
reduce the coupling of seismic motion in the vertical direction to the mirror.
As in the case of the mode cleaner pendulums, collocated feedback systems
are used to damp all six degrees of freedom of the upper pendulum mass and
through cross-coupling the other solid-body modes of the multiple pendulum
system. The control forces for the length and alignment control are applied
from a reaction pendulum which consists of a similar triple pendulum sus-
pended 3mm behind the corresponding mirror. The intermediate mass of the
reaction pendulum carries coils which act on magnets glued to the intermedi-
ate mass of the mirror triple pendulum. To keep the internal quality factor of
the mirrors as high as possible, no magnets are glued to the mirror itself, but
electrostatic feedback between the mirror and the lowest mass of the reaction
pendulum is used to apply feedback forces in the high frequency control band.
To further minimize the mechanical losses and thus internal thermal noise of
the mirrors and the pendulums, the lowest pendulum stage consists entirely
of fused silica; see Figure 6. The mechanical quality factor Q of fused-silica
suspensions comparable in size has been demonstrated to be greater than
2×107. Small fused-silica pieces are attached to the intermediate mass and to
the mirror itself by hydroxide-catalysis bonding. This technique provides
high-strength bonds and allows the high quality factor to be maintained and
therefore the thermal noise to be kept low. Four fused-silica fibers of 270 µm
diameter each are welded to these fused-silica pieces and support the mirrors.
The optical layout of GEO600 (see Figure 7) can be divided into four major
parts: The laser system, the input optics, the dual-recycled Michelson inter-
ferometer, and the output optics followed by the main photodetector. Some
steering mirrors, electro-optical modulators and Faraday isolators are omitted
in Figure 7. All optical components but the laser system and the photodetector
are suspended inside the vacuum system.
Fig. 6. An outline sketch of the test mass suspensions in GEO600. The test mass
and intermediate mass are made of fused-silica and are connected by 4 fused-silica
fibres. The reaction mass is also of fused-silica and has the electrode pattern required
to allow electrostatic actuation forces to be applied to the test mass. The other
masses are fabricated from metal. Two stages of vertical isolation are provided in
the form of cantilever mounted blade springs. Active local damping is provided
from a structure (not shown) held at the upper mass level by the damping arms.
The pendulum chains are suspended from a structure which allows crude angular
alignment of the mirror, and this is in turn supported by 3 vibration isolating legs.
Each leg consists of an active layer and a passive layer which, to avoid contamination
of the vacuum are enclosed in steel bellows. A flex-pivot is then required to provide
rotational compliance. The reaction chain is omitted from the test mass suspensions
in the end stations.
The GEO600 laser system is based on an injection-locked laser-diode
pumped Nd:YAG system with an output power of 12 W. A non-planar ring-
oscillator (NPRO) with an output power of 0.8 W is used as the master laser.
The slave laser is formed by a four mirror cavity with two Nd:YAG rods serv-
ing as gain media. Each crystal is pumped by fiber-coupled laser diodes with
a power of 17W. Two Brewster plates are incorporated in the slave cavity to
define the polarization direction, reduce depolarization losses and compensate
for the astigmatism introduced by the curved mirrors of the slave resonator.
Fig. 7. Optical Layout of GEO600. A 12 W injection locked laser system is filtered by
two sequential mode cleaners and injected into the dual (power and signal) recycled
interferometer. Only power recycling was used for the S1 run. A folded light path
is used to increase the round-trip length of the interferometer arms to 2400 m. An
output mode cleaner will be used to spatially clean the laser mode before it reaches
Due to the mode-selective pumping scheme more than 95% of the light leaving
the laser is in the fundamental TEM00mode. The fully automated injection
locking control servo system acquires lock within 100 ms and allows stable
operation of the laser system.
4.4 Input Optics
The light from the laser system is passed in transmission sequentially through
two triangular resonant cavities, serving as frequency references and optical
mode cleaners; they are of 8.0 m and 8.1 m round-trip lengths. The laser
frequency is stabilized to the resonant frequency of the first mode cleaner
MC1. For this purpose radio frequency phase modulation sidebands are
impressed on the laser beam prior to entering the first mode cleaner. The
light reflected by the input mirror of the first mode cleaner interferes with
the light leaking out of the first mode cleaner on a quadrant photodiode. The
demodulated sum of the photocurrents of all quadrants of this photodiode
is used in the Pound-Drever-Hall scheme to develop an error signal for the
deviation of the laser frequency from a mode cleaner resonance frequency.
This signal is fed back to the master laser frequency actuators and to a phase
correcting Pockels cell to stabilize the laser to the first mode cleaner length.
With this first control loop in place, the laser frequency will change as the
length of MC1 changes. Due to this effect, the length-control actuator of MC1
can be used to bring the laser/MC1 unit into resonance with the second mode
cleaner MC2. For this purpose another pair of rf sidebands is imposed on the
laser beam by an electro-optical modulator, located between the two mode
cleaners. The light reflected by the input mirror of the second mode cleaner
is aligned onto a quadrant photodetector and the sum of the photocurrents of
all segments is demodulated to produce an error signal for this feedback loop.
A third control loop is used to bring the laser/MC1/MC2 unit into resonance
with the power-recycling cavity. A Faraday isolator is used between the mode
cleaner and the power recycling mirror to obtain access to the light reflected
by the power recycling cavity which is detected on a quadrat photodiode. A
detailed description of the frequency control scheme of GEO600 is given in
All the quadrant photodiodes mentioned above are used both for length control
and for wavefront sensing (and thus alignment control) of the mode cleaners
and the power recycling cavity. The difference of the photocurrents of a com-
bination of the quadrants of these diodes is demodulated at the respective
rf frequency and the resulting signals provide alignment information of the
incoming beam relative to the eigenmode of the relevant cavity. A telescope is
used to get near field and far field information of the phase-front differences
which can be converted into tilt or rotation as well as x or y parallel shift in-
formation of the incoming beam relative to the cavity axis z. The appropriate
linear combination of these signals is fed back to the mode cleaner mirrors.
The alignment error signal of the power recycling cavity is used to change the
tilt/rotation of the power recycling mirror and of a steering mirror to keep
this cavity aligned to the incoming beam. The complete automatic alignment
system uses additional quadrant diodes behind several mirrors to keep the
spot positions centered on all the relevant cavity mirrors.
4.5 Interferometer configuration
The main interferometer is designed as a dual-recycled folded-arm Michelson
interferometer[7,8]. Power recycling leads to a power buildup in the interfer-
ometer and improves the shot-noise limited sensitivity of the detector. The
anticipated power buildup in GEO600 is 2000 which results in a power of
about 10kW at the beamsplitter. Any differential phase change of the light
in the interferometer arms (the signature of a gravitational wave) will lead
to a change in the light intensity at the output port of the interferometer.
The partially-transmitting signal recycling mirror will reflect most of this
light back into the interferometer and forms another Fabry-Perot cavity, the
signal-recycling cavity. In this cavity, the light power representing the signal
is enhanced through resonance in a frequency range determined by the cavity
bandwidth of the signal recycling mirror and the resonant frequency of the
signal-recycling cavity. This effect reduces the shot-noise-equivalent apparent
displacement noise of the detector for these frequencies. For the S1 run, the
signal recycling mirror was not yet installed, and so the instrument ran as a
power-recycled folded-arm Michelson.
In the final optical configuration GEO will use an output mode cleaner as a
spatial filter of the main interferometer output beam, placed just before the
antisymmetric photodetector. The output mode cleaner will be installed when
the signal recycling mirror is incorporated, but was not needed for the S1 run.
The length and alignment control systems for the Michelson interferometer
and the signal recycling cavity use similar techniques as described above for
the mode cleaners and power recycling cavity. Quadrant photodiodes sense
the beam at the interferometer output port for the Michelson control. A small
fraction of the light in one interferometer arm is reflected by the anti-reflection
coating of the beamsplitter and is used for the control of the signal recycling
degrees of freedom. The two pairs of sidebands needed for the sensing scheme
are impressed on the laser beam injected into the power recycling cavity, and
the rf frequencies were chosen to be multiples of the free spectral range of the
power recycling cavity. Magnet and coil actuators at the intermediate mass of
the suspensions, and electro-static actuators at the mirror level of the triple
pendulum suspensions, are used as actuators for the length and alignment
GEO600 has five suspended cavities and the suspended Michelson interfer-
ometer which need length and alignment control systems. Thirty pendulums
need local damping of at least 4 degrees of freedom and additional feedback-
control systems are needed for the laser stabilization. Most of these control
loops are implemented with analog electronic controllers with some guidance
by a LabVIEW computer-control environment. Only the active seismic
isolation and some slow alignment-drift-control systems are implemented as
digital control loops. The LabVIEW computer system controls pre-alignment,
guides lock acquisition of the laser and the mode cleaners, monitors the de-
tector status, and compensates for long-term drifts. Typical response times
of this system are 100 ms. The lock acquisition of the recycling cavities and
the interferometer is guided by a micro-controller to allow for faster response
Although only the light at the detector antisymmetric output includes a possi-
ble gravitational-wave signal with a high signal to noise ratio, a multi-channel
data acquisition system is needed to detect environmental and detector dis-
turbances and exclude false detections. Two different sampling rates (16384
Hz and 512 Hz) are used in the data-acquisition system (DAQ) of GEO600.
In the central building 32 fast channels and 64 slow channels are available,
and in each of the end buildings 16 fast channels can be recorded. Most of
these channels will be used for detector characterization only. The data are
recorded into the standard Frame format for later analysis.
5 The S1 Run
The S1 run took place from 23 August 2002 15:00 UTC through 9 September
2002 15:00 UTC. The total duration of 17 days spanned three weekends and
one national holiday in the U.S., which helped reduce the time lost due to
anthropogenic noise sources, particularly at the LIGO Livingston site. Locked
times and duty cycles for the four individual interferometers are given in Table
3, along with the double and triple coincidence times for the LIGO detectors.
The duty cycle of GEO600 (98%) is so high that its coincidence time with
any combination of LIGO interferometers is essentially the same as that of
the LIGO interferometer(s) alone.
Locked times and duty cycles for the S1 Run
Detector/combinationDetector Hours coincidenceLocked/Duty Cycle
LIGO H2 29873.1%
Each LIGO interferometer had a defined operating state, including which ser-
vos were operational, their gains, acceptable light levels on photodiodes, etc.
The instrument operator on duty was responsible to lock the interferometer
and put it into the required configuration, assisted by computer scripts that
set the majority of the parameters. When the desired state was achieved, the
operator issued a command that put the interferometer into “Science mode”,
effectively declaring that the detector was in the proper configuration. At
that time, personnel were restricted from entering the experimental halls. A
computer program began monitoring the computer control network for any
unauthorized changes to the interferometer state, and if any were detected, it
automatically removed the interferometer from Science Mode and raised an
Only data taken in Science Mode segments longer than 300 seconds were
deemed suitable for analysis. However, because of the still incomplete state
of the commissioning, the designation of an interferometer as being in Sci-
ence Mode was not sufficient to ensure that the data were of uniform quality.
Thus each data analysis effort independently evaluated the Science Mode data
(assessing, for example, the noise level or the quality of the calibration) and
made their own selection of data for further analysis, based on the particular
requirements of that analysis.
The LIGO noise level for S1 shown in Figure 8 is substantially above the design
goal. At high frequencies, most of the extra noise can be attributed to the fact
that the interferometers were effectively using very low laser power, and using
the detection sidebands inefficiently, as described above. This leads to a poorer
sensing resolution due to shot noise and electronics noise. At low frequencies,
the excess noise is mainly due to noise in the control systems, typically from
auxiliary control systems such as angular controls. Much of this noise is due
to the incomplete commissioning of the alignment system, and control system
filters which are not yet optimized.
The numerous peaks in the spectrum are due to a number of different sources.
Multiples of 60 Hz are prominent (at the 10−19level in strain), in part due to
switching power supplies that are scheduled for replacement. Acoustic peaks,
due to fans and other rotating mechanical equipment, enter through acoustic
and mechanical coupling mechanisms. Improvements in acoustic shielding and
in the optical layouts to reduce acoustic sensitivity are planned to address
The signal that is analyzed to search for gravitational wave signals is obtained
from the photodiode at the antisymmetric port. Since this is the error signal
for the differential arm length, the effect of the feedback loop gain must be
measured and compensated for. The absolute scale for strain is established
using the laser wavelength, measuring the mirror drive signal required to move
through a given number of interference fringes. The frequency response of the
detector is determined via swept-sine excitations of the end mirrors made
periodically through the run.
One of the difficulties with the S1 data was that drifts in the alignment (be-
cause the wavefront sensing portion of the alignment control system was not
fully operational) caused changes in the coupling of light into the interferome-
Fig. 8. The LIGO interferometer sensitivities for S1. The LHO 4km instrument is
H1; the LHO 2km is H2; and LLO 4km is L1. The ‘SRD Goal’ refers to the LIGO
design sensitivity for the LIGO instruments.
ter. This shows up directly in changes in the optical gain of the interferometers
(e.g., watts per meter at the antisymmetric port), and in changes of the overall
gain of the servo system holding the interferometers in lock. The first effect is
an overall scaling in the signal, and was typically 10% or less in S1. However,
the change in servo gain gives a frequency-dependent correction to the calibra-
tion function, which can be several times larger than the overall scale change,
particularly near the unity gain point of the differential arm servo (between
150 and 200 Hz). To compensate for this problem, the length of one of the
arms was modulated sinusoidally at two frequencies with known amplitudes
using the actuators on one of the end test masses. By monitoring the size of
the resulting signals in the antisymmetric port photodiode signal, the optical
gain of the interferometer could be tracked on a minute-by-minute basis, and
corrections for the drift applied. The calibration procedure and results are
described in more detail in reference ; the overall statistical error is about
Other measures of performance also showed non-stationarity. The band-limited
rms in the antisymmetric port photodiode sometimes showed degradation dur-
Fig. 9. The GEO600 Sensitivity Curve for S1.
100 5001000 2000 5000
Strain (RMS / Hz1/2)
GEO600 noise spectrum (S1 playground 120 av 120 s Hanning window)
ing a locked section as the interferometer drifted away from the initial align-
ment. “Glitch” rates included variations of factors of 3 to 100 between locked
sections. This non-stationary behavior affects some searches more than others,
but should improve as the detectors approach full operation.
The GEO600 sensitivity for the S1 run are shown in Figure 9. The duty cycle
was 98% (see Table 3) and the longest continuous stretch of data is 121h.
During the S1 run GEO600 was operating in the power-recycled Michelson
interferometer configuration with a reduced power recycling gain and reduced
input laser power. The laser power injected into the first mode cleaner was at-
tenuated to 2W. The overall optical transmission of the mode cleaners, phase
modulators and isolators is 52% which leaves approximately 1W of laser light
being injected into the interferometer. The power buildup in the power recy-
cling cavity was 300 which was limited by the 1.3% transmission of the power
recycling mirror installed during S1. The signal recycling mirror was not in-
stalled and test mirrors suspended in wire slings were used for the for the
beamsplitter and the inboard mirrors of the folded interferometer arms.
Due to the automatic alignment system  the longest duration without
manual alignment of the mode cleaners was more than one year prior to S1
and no manual alignment of the mode cleaners was undertaken during the
S1 run. Operator alignment of the power recycling cavity and the Michelson
interferometer was needed only a few times after major seismic disturbances.
An automatic lock acquisition process was initiated whenever any of the cavi-
ties or the Michelson interferometer lost lock. This system and the automatic
alignment system were very stable and reliable and no operator presence at
the detector was required at night.
Calibration of GEO600 during the S1 run was achieved by imposing known
forces on two of the mirrors using electrostatic drives that were fitted to pro-
vide differential control of the interferometer. In the final GEO600 configu-
ration, radiation pressure from a modulated laser beam which will be aligned
onto and reflected by one of the end mirrors will introduce calibration lines
with known amplitude into the output signal of the detector. The spectrum
of the applied calibration force consisted of a line at 244 Hz and its odd
harmonics, generated by suitable filtering from a square wave. The signal gen-
erator used to produce the series was phase-locked to the GPS stabilized clock
to which the data acquisition system was also synchronized. The electrostatic
actuators were calibrated with respect to electro-magnetic actuators one stage
higher up on the same suspension chains. These were in turn calibrated by
applying forces large enough to enable simple fringe counting. The actuators
were found to be adequately linear to allow this calibration method to succeed
with good reliability. The calibrated gravitational wave channel was generated
by measuring the amplitudes of the calibration peaks in 1 s frames of data.
These measurements were used to determine the unknown calibration factors
(essentially the overall optical transfer function magnitude and the gain of
the relevant control loops) by fitting the data to a model based on previously
measured transfer functions of the electro-mechanical control system. The cal-
ibration coefficients were smoothed over periods of one minute. Suitable time
domain digital filters were generated to produce a calibrated gravitational
wave channel. The overall calibration uncertainty was about 4% for signal fre-
quencies above 200 Hz and 6% between 50 Hz and 200 Hz. Further detail of
this process can be found in reference.
The S1 science run is an important milestone for LIGO and GEO, providing
the first data for scientific analysis for two of the newest generation of grav-
itational wave interferometer observatories. Even though the detectors were
operated in a preliminary configuration with many features not implemented,
the data were relatively well behaved, and the sensitivity great enough to
improve on prior observations with broadband gravitational wave detectors.
Several types of analysis have been recently completed using the data from
the S1 run. These include:
• a search for the inspiral signal from binary neutron star mergers
• a search for continuous waves from a rapidly rotating pulsar (J1939+2134)
• a search for short bursts of gravitational waves from unknown sources
• a search for a stochastic background of gravitational waves of cosmological
In all cases, the sensitivity for S1 was not expected to be sufficient to make a
positive detection, so the emphasis in these analyses is to develop techniques
for searching for gravitational waves, to confront the problems of dealing with
real data with its deviations from the usual assumptions of gaussianity and
stationarity, and to set improved upper limits on the flux of gravitational
waves incident on the Earth.
Continued rapid improvements are expected in both the LIGO and GEO de-
tectors. Immediately after the S1 run, GEO commenced the installation and
commissioning of the complete optical configuration by adding the signal recy-
cling mirror. LIGO has continued to complete the control system and tuning,
leading to more than a factor 10 improvement in the sensitivity. New science
data runs are taking place to collect data with this better performance. There
will be a smooth transition from the present epoch, where commissioning dom-
inates, to the goal of effectively continuous astrophysical data collection, as
the instruments approach their goal sensitivity. The data are already inter-
esting in terms of constraining astrophysical models from upper limits, and
significant improvements in the near future will make even better upper limits
possible, along with the increased potential of directly detecting gravitational
waves of astrophysical origin.
The authors gratefully acknowledge the support of the United States National
Science Foundation for the construction and operation of the LIGO Labora-
tory and the Particle Physics and Astronomy Research Council of the United
Kingdom, the Max-Planck-Society and the State of Niedersachsen/Germany
for support of the construction and operation of the GEO600 detector. The
authors also gratefully acknowledge the support of their research by these
agencies and by the Australian Research Council, the Natural Sciences and
Engineering Research Council of Canada, the Council of Scientific and Indus-
trial Research of India, the Department of Science and Technology of India,
the Spanish Ministerio de Ciencia y Tecnologia, the John Simon Guggenheim
Foundation, the David and Lucile Packard Foundation, the Research Corpora-
tion, and the Alfred P. Sloan Foundation. This paper has been assigned LIGO
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