A yellow hypergiant (YHG) is a massive star with an extended atmosphere, a spectral class from A to K, and, starting with an initial mass of about 20–60 solar masses, has lost as much as half that mass. They are amongst the most visually luminous stars, with absolute magnitude (MV) around −9, but also one of the rarest, with just 20 known in the Milky Way and six of those in just a single cluster. They are sometimes referred to as cool hypergiants in comparison with O- and B-type stars, and sometimes as warm hypergiants in comparison with red supergiants.
Classification
The term "hypergiant" was used as early as 1929, but not for the stars currently known as hypergiants. Hypergiants are defined by their '0' luminosity class, and are higher in luminosity than the brightest supergiants of class Ia, although they were not referred to as hypergiants until the late 1970s. Another criterion for hypergiants was also suggested in 1979 for some other highly luminous mass-losing hot stars, but was not applied to cooler stars. In 1991, Rho Cassiopeiae was the first to be described as a yellow hypergiant, likely becoming grouped as a new class of luminous stars during discussions at the Solar physics and astrophysics at interferometric resolution workshop in 1992.
Definitions of the term hypergiant remain vague, and although
luminosity class 0 is for hypergiants, they are more commonly designated
by the alternative luminosity classes Ia-0 and Ia+.
Their great stellar luminosities are determined from various spectral
features, which are sensitive to surface gravity, such as Hβ line widths
in hot stars or a strong Balmer discontinuity in cooler stars. Lower surface gravity often indicates larger stars, and hence, higher luminosities.
In cooler stars, the strength of observed oxygen lines, such as O I at
777.4 nm., can be used to calibrate directly against stellar
luminosity.
One astrophysical method used to definitively identify yellow hypergiants is the so-called Keenan-Smolinski criterion. Here all absorption lines should be strongly broadened, beyond those expected of bright supergiant stars, and also show strong evidence of significant mass loss. Furthermore, at least one broadened Hα
component should also be present. They may also display very complex Hα
profiles, typically having strong emission lines combined with
absorption lines.
The terminology of yellow hypergiants is further complicated by
referring to them as either cool hypergiants or warm hypergiants,
depending on the context. Cool hypergiants refers to all sufficiently
luminous and unstable stars cooler than blue hypergiants and LBVs, including both yellow and red hypergiants. The term warm hypergiants has been used for highly luminous class A and F stars in M31 and M33 that are not LBVs, as well as more generally for yellow hypergiants.
Characteristics
Visual light curve for ρ Cassiopeiae from 1933 to 2015
Yellow hypergiants occupy a region of the Hertzsprung–Russell diagram above the instability strip,
a region where relatively few stars are found and where those stars are
generally unstable. The spectral and temperature ranges are
approximately A0-K2 and 4,000–8,000 K (3,730–7,730 °C; 6,740–13,940 °F)
respectively. The area is bounded on the high-temperature side by the Yellow Evolutionary Void
where stars of this luminosity become extremely unstable and experience
severe mass loss. The “Yellow Evolutionary Void” separates yellow
hypergiants from luminous blue variables
although yellow hypergiants at their hottest and luminous blue
variables at their coolest can have approximately the same temperature
near 8,000 K. At the lower temperature bound, yellow hypergiants and red
supergiants are not clearly separated; RW Cephei (roughly 4,000 K (3,730 °C; 6,740 °F), 295,000 L☉) is an example of a star that shares characteristics of both yellow hypergiants and red supergiants.
Yellow hypergiants have a fairly narrow range of luminosities above 200,000 L☉ (e.g. V382 Carinae at 212,000 L☉) and below the Humphrey-Davidson limit at around 600,000 L☉.
With their output peaking in the middle of the visual range, these are
the most visually bright stars known with absolute magnitudes around −9
or −9.5 .
They are large and somewhat unstable, with very low surface gravities. Where yellow supergiants
have surface gravities (log g) below about 2, the yellow hypergiants
have log g around zero. In addition they pulsate irregularly, producing
small variations in temperature and brightness. This produces very
high mass loss rates, and nebulosity is common around the stars. Occasional larger outbursts can temporarily obscure the stars.
Yellow hypergiants form from massive stars after they have
evolved away from the main sequence. Most observed yellow hypergiants
have been through a red supergiant phase and are evolving back towards
higher temperatures, but a few are seen in the brief first transition
from main sequence to red supergiant. Supergiants with an initial mass
less than 20 M☉ will explode as a supernova while still red supergiants, while stars more massive than about 60 M☉ will never cool beyond blue supergiant temperatures. The exact mass ranges depend on metallicity and rotation. Yellow supergiants cooling for the first time may be massive stars of up to 60 M☉ or more, but post-red supergiant stars will have lost around half their initial mass.
Chemically, most yellow hypergiants show strong surface enhancement of nitrogen and also of sodium and some other heavy elements. Carbon and oxygen are depleted, while helium is enhanced, as expected for a post-main-sequence star.
Evolution
Yellow
hypergiants have clearly evolved off the main sequence and so have
depleted the hydrogen in their cores. The majority of yellow
hypergiants are postulated to be post-red supergiants evolving blueward,
while more stable and less luminous yellow supergiants are likely to be
evolving to red supergiants for the first time. There is strong
chemical and surface gravity evidence that the brightest of the yellow
supergiants, HD 33579, is currently expanding from a blue supergiant to a red supergiant.
These stars are doubly rare because they are very massive,
initially hot class O-type main-sequence stars more than 15 times as
massive as the Sun, but also because they spend only a few thousand
years in the unstable yellow void phase of their lives. In fact, it is
difficult to explain even the small number of observed yellow
hypergiants, relative to red supergiants of comparable luminosity, from
simple models of stellar evolution. The most luminous red supergiants
may execute multiple "blue loops", shedding much of their atmosphere,
but without actually ever reaching the blue supergiant stage, each one
taking only a few decades at most. Conversely, some apparent yellow
hypergiants may be hotter stars, such as the "missing" LBVs, masked
within a cool pseudo-photosphere.
Recent discoveries of blue supergiant supernova progenitors have
also raised the question of whether stars could explode directly from
the yellow hypergiant stage.
A handful of possible yellow supergiant supernova progenitors have
been discovered, but they all appear to be of relatively low mass and
luminosity, not hypergiants. SN 2013cu
is a type IIb supernova whose progenitor has been directly and clearly
observed. It was an evolved star around 8,000 K (7,730 °C; 13,940 °F)
showing extreme mass loss of helium and nitrogen enriched material.
Although the luminosity is not known, only a yellow hypergiant or
luminous blue variable in outburst would have these properties.
Modern models suggest that stars with a certain range of masses and rotation rates may explode as supernovae
without ever becoming blue supergiants again, but many will eventually
pass right through the yellow void and become low-mass low-luminosity luminous blue variables and possibly Wolf–Rayet stars after that.
Specifically, more massive stars and those with higher mass loss rates
due to rotation or high metallicity will evolve beyond the yellow
hypergiant stage to hotter temperatures before reaching core collapse.
Structure
IRAS 17163-3907 is a yellow hypergiant that clearly shows the expelled material that probably surrounds all yellow hypergiants.
According to the current physical models of stars, a yellow hypergiant should possess a convective core surrounded by a radiative zone, as opposed to a sun-sized star, which consists of a radiative core surrounded by a convective zone. Because of their extreme luminosity and internal structure, yellow hypergiants suffer high rates of mass loss and are generally surrounded by envelopes of expelled material. An example of the nebulae that can result is IRAS 17163-3907, known as the Fried Egg, which has expelled several solar masses of material in just a few hundred years.
The yellow hypergiant is an expected phase of evolution as the
most luminous red supergiants evolve bluewards, but they may also
represent a different sort of star. LBVs during eruption have such
dense winds that they form a pseudo-photosphere which appears as a
larger cooler star despite the underlying blue supergiant being largely
unchanged. These are observed to have a very narrow range of
temperatures around 8,000 K (7,730 °C; 13,940 °F). At the bistability
jump which occurs around 21,000 K (20,700 °C; 37,300 °F) blue supergiant
winds become several times denser and could be result in an even cooler
pseudo-photosphere. No LBVs are observed just below the luminosity
where the bistability jump crosses the S Doradus instability strip (not to be confused with the Cepheid instability strip), but it is theorised that they do exist and appear as yellow hypergiants because of their pseudo-photospheres.
Known yellow hypergiants
Yellow hypergiant HR 5171 A, seen as the bright yellow star at the center of the image.Artist's impression of the binary system containing yellow hypergiant HR 5171 A
The Laser Interferometer Gravitational-Wave Observatory (LIGO) is a large-scale physics experiment and observatory designed to detect cosmic gravitational waves and to develop gravitational-wave observations as an astronomical tool. Two large observatories were built in the United States with the aim of detecting gravitational waves by laserinterferometry.
These observatories use mirrors spaced four kilometers apart which are
capable of detecting a change of less than one ten-thousandth the charge diameter of a proton.
The initial LIGO observatories were funded by the United States National Science Foundation (NSF) and were conceived, built and are operated by Caltech and MIT. They collected data from 2002 to 2010 but no gravitational waves were detected.
The Advanced LIGO Project to enhance the original LIGO detectors
began in 2008 and continues to be supported by the NSF, with important
contributions from the United Kingdom's Science and Technology Facilities Council, the Max Planck Society of Germany, and the Australian Research Council. The improved detectors began operation in 2015. The detection of gravitational waves was reported in 2016 by the LIGO Scientific Collaboration (LSC) and the Virgo Collaboration
with the international participation of scientists from several
universities and research institutions. Scientists involved in the
project and the analysis of the data for gravitational-wave astronomy are organized by the LSC, which includes more than 1000 scientists worldwide, as well as 440,000 active Einstein@Home users as of December 2016.
LIGO is the largest and most ambitious project ever funded by the NSF. In 2017, the Nobel Prize in Physics was awarded to Rainer Weiss, Kip Thorne and Barry C. Barish "for decisive contributions to the LIGO detector and the observation of gravitational waves".
Observations are made in "runs". As of January 2022, LIGO has made three runs (with one of the runs divided into two "subruns"), and made 90 detections of gravitational waves.
Maintenance and upgrades of the detectors are made between runs. The
first run, O1, which ran from 12 September 2015 to 19 January 2016, made
the first three detections, all black hole mergers. The second run, O2,
which ran from 30 November 2016 to 25 August 2017, made eight
detections, seven black hole mergers, and the first neutron star merger.
The third run, O3, began on 1 April 2019; it was divided into O3a, from
1 April to 30 September 2019, and O3b, from 1 November 2019 until it was suspended on 27 March 2020 due to COVID-19. The O3 run made the first detection of a merger of a neutron star with a black hole.
The gravitational wave observatories LIGO, Virgo, and KAGRA are coordinating to continue observations after the COVID-caused stop, and LIGO's O4 observing run started on May 24, 2023. LIGO projects a sensitivity goal of 160–190 Mpc for binary neutron star
mergers (sensitivities: Virgo 80–115 Mpc, KAGRA greater than 1 Mpc).
History
Background
LIGO Hanford ObservatoryLIGO
Laboratory operates two detector sites, one near Hanford in eastern
Washington, and another near Livingston, Louisiana. This photo shows the
Livingston detector site.
The LIGO concept built upon early work by many scientists to test a component of Albert Einstein's theory of general relativity, the existence of gravitational waves. Starting in the 1960s, American scientists including Joseph Weber, as well as Soviet scientists Mikhail Gertsenshtein and Vladislav Pustovoit, conceived of basic ideas and prototypes of laser interferometry, and in 1967 Rainer Weiss of MIT
published an analysis of interferometer use and initiated the
construction of a prototype with military funding, but it was terminated
before it could become operational. Starting in 1968, Kip Thorne initiated theoretical efforts on gravitational waves and their sources at Caltech, and was convinced that gravitational wave detection would eventually succeed.
Prototype interferometric gravitational wave detectors (interferometers) were built in the late 1960s by Robert L. Forward and colleagues at Hughes Research Laboratories
(with mirrors mounted on a vibration isolated plate rather than free
swinging), and in the 1970s (with free swinging mirrors between which
light bounced many times) by Weiss at MIT, and then by Heinz Billing and colleagues in Garching Germany, and then by Ronald Drever, James Hough and colleagues in Glasgow, Scotland.
In 1980, the NSF funded the study of a large interferometer led by MIT (Paul Linsay, Peter Saulson,
Rainer Weiss), and the following year, Caltech constructed a 40-meter
prototype (Ronald Drever and Stan Whitcomb). The MIT study established
the feasibility of interferometers at a 1-kilometer scale with adequate
sensitivity.
Under pressure from the NSF, MIT and Caltech were asked to join
forces to lead a LIGO project based on the MIT study and on experimental
work at Caltech, MIT, Glasgow, and Garching.
Drever, Thorne, and Weiss formed a LIGO steering committee, though they
were turned down for funding in 1984 and 1985. By 1986, they were asked
to disband the steering committee and a single director, Rochus E. Vogt (Caltech), was appointed. In 1988, a research and development proposal achieved funding.
From 1989 through 1994, LIGO failed to progress technically and
organizationally. Only political efforts continued to acquire funding. Ongoing funding was routinely rejected until 1991, when the U.S. Congress
agreed to fund LIGO for the first year for $23 million. However,
requirements for receiving the funding were not met or approved, and the
NSF questioned the technological and organizational basis of the
project. By 1992, LIGO was restructured with Drever no longer a direct participant. Ongoing project management issues and technical concerns were revealed
in NSF reviews of the project, resulting in the withholding of funds
until they formally froze spending in 1993.
In 1994, after consultation between relevant NSF personnel,
LIGO's scientific leaders, and the presidents of MIT and Caltech, Vogt
stepped down and Barry Barish (Caltech) was appointed laboratory director, and the NSF made clear that LIGO had one last chance for support.
Barish's team created a new study, budget, and project plan with a
budget exceeding the previous proposals by 40%. Barish proposed to the
NSF and National Science Board to build LIGO as an evolutionary
detector, where detection of gravitational waves with initial LIGO would
be possible, and with advanced LIGO would be probable.This new proposal received NSF funding, Barish was appointed Principal Investigator,
and the increase was approved. In 1994, with a budget of US$395
million, LIGO stood as the largest overall funded NSF project in
history. The project broke ground in Hanford, Washington in late 1994
and in Livingston, Louisiana in 1995. As construction neared completion
in 1997, under Barish's leadership two organizational institutions were
formed, LIGO Laboratory and LIGO Scientific Collaboration (LSC). The
LIGO laboratory consists of the facilities supported by the NSF under
LIGO Operation and Advanced R&D; this includes administration of the
LIGO detector and test facilities. The LIGO Scientific Collaboration is
a forum for organizing technical and scientific research in LIGO. It is
a separate organization from LIGO Laboratory with its own oversight.
Barish appointed Weiss as the first spokesperson for this scientific
collaboration.
Observations begin
Initial
LIGO operations between 2002 and 2010 did not detect any gravitational
waves. In 2004, under Barish, the funding and groundwork were laid for
the next phase of LIGO development (called "Enhanced LIGO"). This was
followed by a multi-year shut-down while the detectors were replaced by
much improved "Advanced LIGO" versions. Much of the research and development work for the LIGO/aLIGO machines was based on pioneering work for the GEO600 detector at Hannover, Germany. By February 2015, the detectors were brought into engineering mode in both locations.
By mid-September 2015, "the world's largest gravitational-wave
facility" completed a 5-year US$200-million overhaul at a total cost of
$620 million.
On 18 September 2015, Advanced LIGO began its first formal science
observations at about four times the sensitivity of the initial LIGO
interferometers. Its sensitivity was to be further enhanced until it was planned to reach design sensitivity around 2021.
Current executive director David Reitze
announced the findings at a media event in Washington D.C., while
executive director emeritus Barry Barish presented the first scientific
paper of the findings at CERN to the physics community.
On 16 June 2016 LIGO announced a second signal
was detected from the merging of two black holes with 14.2 and 7.5
times the mass of the Sun. The signal was picked up on 26 December 2015,
at 3:38 UTC.
The detection of a third black hole merger, between objects of
31.2 and 19.4 solar masses, occurred on 4 January 2017 and was announced
on 1 June 2017. Laura Cadonati was appointed the first deputy spokesperson.
A fourth detection of a black hole merger, between objects of
30.5 and 25.3 solar masses, was observed on 14 August 2017 and was
announced on 27 September 2017.
In 2017, Weiss, Barish, and Thorne received the Nobel Prize in Physics
"for decisive contributions to the LIGO detector and the observation of
gravitational waves." Weiss was awarded one-half of the total prize
money, and Barish and Thorne each received a one-quarter prize.
After shutting down for improvements, LIGO resumed operation on
26 March 2019, with Virgo joining the network of gravitational-wave
detectors on 1 April 2019. Both ran until 27 March 2020, when the COVID-19 pandemic halted operations.
During the COVID shutdown, LIGO underwent a further upgrade in
sensitivity, and observing run O4 with the new sensitivity began on 24
May 2023.
Mission
Detector
noise curves for Initial and Advanced LIGO as a function of frequency.
They lie above the bands for space-borne detectors like the evolved Laser Interferometer Space Antenna (eLISA) and pulsar timing arrays such as the European Pulsar Timing Array
(EPTA). The characteristic strains of potential astrophysical sources
are also shown. To be detectable the characteristic strain of a signal
must be above the noise curve. These frequencies that aLIGO can detect are in the range of human hearing.
LIGO's mission is to directly observe gravitational waves of cosmic origin. These waves were first predicted by Einstein's general theory of relativity
in 1916, when the technology necessary for their detection did not yet
exist. Their existence was indirectly confirmed when observations of the
binary pulsar PSR 1913+16 in 1974 showed an orbital decay which matched Einstein's predictions of energy loss by gravitational radiation. The Nobel Prize in Physics 1993 was awarded to Hulse and Taylor for this discovery.
Direct detection of gravitational waves had long been sought.
Their discovery has launched a new branch of astronomy to complement electromagnetic telescopes and neutrino observatories. Joseph Weber pioneered the effort to detect gravitational waves in the 1960s through his work on resonant mass bar detectors. Bar detectors continue to be used at six sites worldwide. By the 1970s, scientists including Rainer Weiss realized the applicability of laser interferometry to gravitational wave measurements. Robert Forward operated an interferometric detector at Hughes in the early 1970s.
In fact as early as the 1960s, and perhaps before that, there
were papers published on wave resonance of light and gravitational
waves. Work was published in 1971 on methods to exploit this resonance for the detection of high-frequency gravitational waves.
In 1962, M. E. Gertsenshtein and V. I. Pustovoit published the very
first paper describing the principles for using interferometers for the
detection of very long wavelength gravitational waves. The authors argued that by using interferometers the sensitivity can be 107 to 1010 times better than by using electromechanical experiments. Later, in 1965, Braginsky
extensively discussed gravitational-wave sources and their possible
detection. He pointed out the 1962 paper and mentioned the possibility
of detecting gravitational waves if the interferometric technology and
measuring techniques improved.
Since the early 1990s, physicists have thought that technology has evolved to the point where detection of gravitational waves—of significant astrophysical interest—is now possible.
In August 2002, LIGO began its search for cosmic gravitational
waves. Measurable emissions of gravitational waves are expected from
binary systems (collisions and coalescences of neutron stars or black holes), supernova
explosions of massive stars (which form neutron stars and black holes),
accreting neutron stars, rotations of neutron stars with deformed
crusts, and the remnants of gravitational radiation created by the birth of the universe.
The observatory may, in theory, also observe more exotic hypothetical
phenomena, such as gravitational waves caused by oscillating cosmic strings or colliding domain walls.
Observatories
LIGO operates two gravitational wave observatories in unison: the LIGO Livingston Observatory (30°33′46.42″N90°46′27.27″W) in Livingston, Louisiana, and the LIGO Hanford Observatory, on the DOE Hanford Site (46°27′18.52″N119°24′27.56″W), located near Richland, Washington.
These sites are separated by 3,002 kilometers (1,865 miles) straight
line distance through the earth, but 3,030 kilometers (1,883 miles) over
the surface. Since gravitational waves are expected to travel at the
speed of light, this distance corresponds to a difference in
gravitational wave arrival times of up to ten milliseconds. Through the
use of trilateration, the difference in arrival times helps to determine the source of the wave, especially when a third similar instrument like Virgo, located at an even greater distance in Europe, is added.
Each observatory supports an L-shaped ultra high vacuum system, measuring four kilometers (2.5 miles) on each side. Up to five interferometers can be set up in each vacuum system.
The LIGO Livingston Observatory houses one laser interferometer
in the primary configuration. This interferometer was successfully
upgraded in 2004 with an active vibration isolation system based on
hydraulic actuators providing a factor of 10 isolation in the 0.1–5 Hz
band. Seismic vibration in this band is chiefly due to microseismic waves and anthropogenic sources (traffic, logging, etc.).
The LIGO Hanford Observatory houses one interferometer, almost
identical to the one at the Livingston Observatory. During the Initial
and Enhanced LIGO phases, a half-length interferometer operated in
parallel with the main interferometer. For this 2 km interferometer, the
Fabry–Pérot
arm cavities had the same optical finesse, and, thus, half the storage
time as the 4 km interferometers. With half the storage time, the
theoretical strain sensitivity was as good as the full length
interferometers above 200 Hz but only half as good at low frequencies.
During the same era, Hanford retained its original passive seismic
isolation system due to limited geologic activity in Southeastern
Washington.
Operation
Simplified operation of a gravitational wave observatory
Figure 1:
A beamsplitter (green line) splits coherent light (from the white box)
into two beams which reflect off the mirrors (cyan oblongs); only one
outgoing and reflected beam in each arm is shown, and separated for
clarity. The reflected beams recombine and an interference pattern is
detected (purple circle).
Figure 2: A gravitational wave passing over the left arm (yellow) changes its length and thus the interference pattern.
The parameters in this section refer to the Advanced LIGO experiment. The primary interferometer consists of two beam lines of 4 km length which form a power-recycled Michelson interferometer with Gires–Tournois etalon arms. A pre-stabilized 1064 nm Nd:YAG laser
emits a beam with a power of 20 W that passes through a power recycling
mirror. The mirror fully transmits light incident from the laser and
reflects light from the other side increasing the power of the light
field between the mirror and the subsequent beam splitter to 700 W. From
the beam splitter the light travels along two orthogonal arms. By the
use of partially reflecting mirrors, Fabry–Pérot cavities are created in both arms that increase the effective path length of laser light in the arm from 4 km to approximately 1,200 km. The power of the light field in the cavity is 100 kW.
When a gravitational wave passes through the interferometer, the
spacetime in the local area is altered. Depending on the source of the
wave and its polarization, this results in an effective change in length
of one or both of the cavities. The effective length change between the
beams will cause the light currently in the cavity to become very
slightly out of phase (antiphase) with the incoming light. The cavity will therefore periodically get very slightly out of coherence and the beams, which are tuned to destructively interfere at the detector, will have a very slight periodically varying detuning. This results in a measurable signal.
After an equivalent of approximately 280 trips down the 4 km length to the far mirrors and back again,
the two separate beams leave the arms and recombine at the beam
splitter. The beams returning from two arms are kept out of phase so
that when the arms are both in coherence and interference (as when there
is no gravitational wave passing through), their light waves subtract,
and no light should arrive at the photodiode.
When a gravitational wave passes through the interferometer, the
distances along the arms of the interferometer are shortened and
lengthened, causing the beams to become slightly less out of phase. This
results in the beams coming in phase, creating a resonance,
hence some light arrives at the photodiode and indicates a signal.
Light that does not contain a signal is returned to the interferometer
using a power recycling mirror, thus increasing the power of the light
in the arms.
In actual operation, noise sources can cause movement in the
optics, producing similar effects to real gravitational wave signals; a
great deal of the art and complexity in the instrument is in finding
ways to reduce these spurious motions of the mirrors. Background noise and unknown errors (which happen daily) are in the order of 10−20, while gravitational wave signals are around 10−22. After noise reduction, a signal-to-noise ratio around 20 can be achieved, or higher when combined with other gravitational wave detectors around the world.
Based on current models of astronomical events, and the predictions of the general theory of relativity,
gravitational waves that originate tens of millions of light years from
Earth are expected to distort the 4-kilometre (2.5 mi) mirror spacing
by about 10−18 m, less than one-thousandth the charge diameter of a proton. Equivalently, this is a relative change in distance of approximately one part in 1021. A typical event which might cause a detection event would be the late stage inspiral and merger of two 10-solar-mass
black holes, not necessarily located in the Milky Way galaxy, which is
expected to result in a very specific sequence of signals often
summarized by the slogan chirp,burst,quasi-normal mode ringing,exponential decay.
In their fourth Science Run at the end of 2004, the LIGO
detectors demonstrated sensitivities in measuring these displacements to
within a factor of two of their design.
During LIGO's fifth Science Run in November 2005, sensitivity
reached the primary design specification of a detectable strain of one
part in 1021 over a 100 Hz
bandwidth. The baseline inspiral of two roughly solar-mass neutron
stars is typically expected to be observable if it occurs within about
8 million parsecs (26×106ly), or the vicinity of the Local Group, averaged over all directions and polarizations. Also at this time, LIGO and GEO 600 (the German-UK interferometric detector) began a joint science run, during which they collected data for several months. Virgo
(the French-Italian interferometric detector) joined in May 2007. The
fifth science run ended in 2007, after extensive analysis of data from
this run did not uncover any unambiguous detection events.
In February 2007, GRB 070201, a short gamma-ray burst arrived at Earth from the direction of the Andromeda Galaxy.
The prevailing explanation of most short gamma-ray bursts is the
merger of a neutron star with either a neutron star or a black hole.
LIGO reported a non-detection for GRB 070201, ruling out a merger at the
distance of Andromeda with high confidence. Such a constraint was
predicated on LIGO eventually demonstrating a direct detection of
gravitational waves.
After the completion of Science Run 5, initial LIGO was upgraded with
certain technologies, planned for Advanced LIGO but available and able
to be retrofitted to initial LIGO, which resulted in an
improved-performance configuration dubbed Enhanced LIGO. Some of the improvements in Enhanced LIGO included:
Science Run 6 (S6) began in July 2009 with the enhanced configurations on the 4 km detectors. It concluded in October 2010, and the disassembly of the original detectors began.
Advanced LIGO
Simplified diagram of an Advanced LIGO detector (not to scale).Design sensitivity of Advanced LIGO interferometer with major noise sources, maximum sensitivity is around 500 Hz
After 2010, LIGO went offline for several years for a major upgrade,
installing the new Advanced LIGO detectors in the LIGO Observatory
infrastructures.
The project continued to attract new members, with the Australian National University and University of Adelaide
contributing to Advanced LIGO, and by the time the LIGO Laboratory
started the first observing run 'O1' with the Advanced LIGO detectors in
September 2015, the LIGO Scientific Collaboration included more than
900 scientists worldwide.
The first observing run operated at a sensitivity roughly three times greater than Initial LIGO, and a much greater sensitivity for larger systems with their peak radiation at lower audio frequencies.
On 11 February 2016, the LIGO and Virgo collaborations announced the first observation of gravitational waves. The signal was named GW150914.
The waveform showed up on 14 September 2015, within just two days of
when the Advanced LIGO detectors started collecting data after their
upgrade. It matched the predictions of general relativity for the inward spiral and merger of a pair of black holes
and subsequent ringdown of the resulting single black hole. The
observations demonstrated the existence of binary stellar-mass black
hole systems and the first observation of a binary black hole merger.
On 15 June 2016, LIGO announced the detection of a second
gravitational wave event, recorded on 26 December 2015, at 3:38 UTC.
Analysis of the observed signal indicated that the event was caused by
the merger of two black holes with masses of 14.2 and 7.5 solar masses,
at a distance of 1.4 billion light years. The signal was named GW151226.
The second observing run (O2) ran from 30 November 2016 to 25 August 2017, with Livingston achieving 15–25% sensitivity improvement over O1, and with Hanford's sensitivity similar to O1. In this period, LIGO saw several further gravitational wave events: GW170104 in January; GW170608 in June; and five others between July and August 2017. Several of these were also detected by the Virgo Collaboration. Unlike the black hole mergers which are only detectable gravitationally, GW170817 came from the collision of two neutron stars and was also detected electromagnetically by gamma ray satellites and optical telescopes.
The third run (O3) began on 1 April 2019 and was planned to last until 30 April 2020; in fact it was suspended in March 2020 due to COVID-19.
On 6 January 2020, LIGO announced the detection of what appeared to be
gravitational ripples from a collision of two neutron stars, recorded on
25 April 2019, by the LIGO Livingston detector. Unlike GW170817, this
event did not result in any light being detected. Furthermore, this is
the first published event for a single-observatory detection, given that
the LIGO Hanford detector was temporarily offline at the time and the
event was too faint to be visible in Virgo's data.
Future observing runs will be interleaved with commissioning
efforts to further improve the sensitivity. It was aimed to achieve
design sensitivity in 2021; the next observing run (O4) was planned to start in December 2022, but the date was pushed back to 24 May 2023.
LIGO-India,
or INDIGO, is a planned collaborative project between the LIGO
Laboratory and the Indian Initiative in Gravitational-wave Observations
(IndIGO) to create a gravitational-wave detector in India. The LIGO
Laboratory, in collaboration with the US National Science Foundation
and Advanced LIGO partners from the U.K., Germany and Australia, has
offered to provide all of the designs and hardware for one of the three
planned Advanced LIGO detectors to be installed, commissioned, and
operated by an Indian team of scientists in a facility to be built in
India.
The LIGO-India project is a collaboration between LIGO Laboratory
and the LIGO-India consortium: Institute of Plasma Research,
Gandhinagar; IUCAA (Inter-University Centre for Astronomy and
Astrophysics), Pune and Raja Ramanna Centre for Advanced Technology,
Indore.
The expansion of worldwide activities in gravitational-wave
detection to produce an effective global network has been a goal of LIGO
for many years. In 2010, a developmental roadmap issued by the Gravitational Wave International Committee (GWIC)
recommended that an expansion of the global array of interferometric
detectors be pursued as a highest priority. Such a network would afford
astrophysicists with more robust search capabilities and higher
scientific yields. The current agreement between the LIGO Scientific
Collaboration and the Virgo collaboration links three detectors of
comparable sensitivity and forms the core of this international network.
Studies indicate that the localization of sources by a network that
includes a detector in India would provide significant improvements.
Improvements in localization averages are predicted to be approximately
an order of magnitude, with substantially larger improvements in
certain regions of the sky.
The NSF
was willing to permit this relocation, and its consequent schedule
delays, as long as it did not increase the LIGO budget. Thus, all costs
required to build a laboratory equivalent to the LIGO sites to house
the detector would have to be borne by the host country. The first potential distant location was at AIGO in Western Australia, however the Australian government was unwilling to commit funding by 1 October 2011 deadline.
A location in India was discussed at a Joint Commission meeting between India and the US in June 2012.
In parallel, the proposal was evaluated by LIGO's funding agency, the
NSF. As the basis of the LIGO-India project entails the transfer of one
of LIGO's detectors to India, the plan would affect work and scheduling
on the Advanced LIGO upgrades already underway. In August 2012, the
U.S. National Science Board approved the LIGO Laboratory's request to
modify the scope of Advanced LIGO by not installing the Hanford "H2"
interferometer, and to prepare it instead for storage in anticipation of
sending it to LIGO-India. In India, the project was presented to the Department of Atomic Energy and the Department of Science and Technology
for approval and funding. On 17 February 2016, less than a week after
LIGO's landmark announcement about the detection of gravitational waves,
Indian Prime Minister Narendra Modi announced that the Cabinet has granted 'in-principle' approval to the LIGO-India mega science proposal.
On 7 April 2023, the LIGO-India project was approved by the
Cabinet of Government of India. Construction is to begin in
Maharashtra's Hingoli district at a cost of INR 2600 crores.
A+
Like Enhanced LIGO, certain improvements will be retrofitted to the existing Advanced LIGO instrument. These are referred to as A+ proposals, and are planned for installation starting from 2019 until the upgraded detector is operational in 2024. The changes would almost double Advanced LIGO's sensitivity, and increase the volume of space searched by a factor of seven. The upgrades include:
Because the final LIGO output photodetector is sensitive to phase,
and not amplitude, it is possible to squeeze the signal so there is less
phase noise and more amplitude noise, without violating the quantum mechanical limit on their product.
This is done by injecting a "squeezed vacuum state" into the dark port
(interferometer output) which is quieter, in the relevant parameter,
than simple darkness. Such a squeezing upgrade was installed at both
LIGO sites prior to the third observing run. The A+ improvement will see the installation of an additional optical cavity
that acts to rotate the squeezing quadrature from phase-squeezed at
high frequencies (above 50 Hz) to amplitude-squeezed at low frequencies,
thereby also mitigating low-frequency radiation pressure noise.
LIGO Voyager
A
third-generation detector at the existing LIGO sites is being planned
under the name "LIGO Voyager" to improve the sensitivity by an
additional factor of two, and halve the low-frequency cutoff to 10 Hz. Plans call for the glass mirrors and 1064 nm lasers to be replaced by even larger 160 kg silicon test masses, cooled to 123 K (a temperature achievable with liquid nitrogen),
and a change to a longer laser wavelength in the 1500–2200 nm range at
which silicon is transparent. (Many documents assume a wavelength of
1550 nm, but this is not final.)
Voyager would be an upgrade to A+, to be operational around 2027–2028.
Cosmic Explorer
A design for a larger facility with longer arms is called "Cosmic Explorer".
This is based on the LIGO Voyager technology, has a similar LIGO-type
L-shape geometry but with 40 km arms. The facility is currently planned
to be on the surface. It has a higher sensitivity than Einstein Telescope for frequencies beyond 10 Hz, but lower sensitivity under 10 Hz.
There are several methods for measuring the temporal coherence of the light (see: field-autocorrelation), including the continuous-wave and the pulsed Fourier-transform spectrometer or Fourier-transform spectrograph.
The term "Fourier-transform spectroscopy" reflects the fact that in all these techniques, a Fourier transform is required to turn the raw data into the actual spectrum, and in many of the cases in optics involving interferometers, is based on the Wiener–Khinchin theorem.
Conceptual introduction
Measuring an emission spectrum
An example of a spectrum: The spectrum of light emitted by the blue flame of a butane torch. The horizontal axis is the wavelength of light, and the vertical axis represents how much light is emitted by the torch at that wavelength.
One of the most basic tasks in spectroscopy is to characterize the spectrum
of a light source: how much light is emitted at each different
wavelength. The most straightforward way to measure a spectrum is to
pass the light through a monochromator, an instrument that blocks all of the light except
the light at a certain wavelength (the un-blocked wavelength is set by a
knob on the monochromator). Then the intensity of this remaining
(single-wavelength) light is measured. The measured intensity directly
indicates how much light is emitted at that wavelength. By varying the
monochromator's wavelength setting, the full spectrum can be measured.
This simple scheme in fact describes how some spectrometers work.
Fourier-transform spectroscopy is a less intuitive way to get the
same information. Rather than allowing only one wavelength at a time to
pass through to the detector, this technique lets through a beam
containing many different wavelengths of light at once, and measures the
total beam intensity. Next, the beam is modified to contain a different
combination of wavelengths, giving a second data point. This process is
repeated many times. Afterwards, a computer takes all this data and
works backwards to infer how much light there is at each wavelength.
To be more specific, between the light source and the detector,
there is a certain configuration of mirrors that allows some wavelengths
to pass through but blocks others (due to wave interference).
The beam is modified for each new data point by moving one of the
mirrors; this changes the set of wavelengths that can pass through.
As mentioned, computer processing is required to turn the raw
data (light intensity for each mirror position) into the desired result
(light intensity for each wavelength). The processing required turns out
to be a common algorithm called the Fourier transform
(hence the name, "Fourier-transform spectroscopy"). The raw data is
sometimes called an "interferogram". Because of the existing computer
equipment requirements, and the ability of light to analyze very small
amounts of substance, it is often beneficial to automate many aspects of
the sample preparation. The sample can be better preserved and the
results are much easier to replicate. Both of these benefits are
important, for instance, in testing situations that may later involve
legal action, such as those involving drug specimens.
Measuring an absorption spectrum
An "interferogram" from a Fourier-transform spectrometer. This is the "raw data" which can be Fourier-transformed
into an actual spectrum. The peak at the center is the ZPD position
("zero path difference"): Here, all the light passes through the interferometer because its two arms have equal length.
The method of Fourier-transform spectroscopy can also be used for absorption spectroscopy. The primary example is "FTIR Spectroscopy", a common technique in chemistry.
In general, the goal of absorption spectroscopy is to measure how
well a sample absorbs or transmits light at each different wavelength.
Although absorption spectroscopy and emission spectroscopy are different
in principle, they are closely related in practice; any technique for
emission spectroscopy can also be used for absorption spectroscopy.
First, the emission spectrum of a broadband lamp is measured (this is
called the "background spectrum"). Second, the emission spectrum of the
same lamp shining through the sample is measured (this is called
the "sample spectrum"). The sample will absorb some of the light,
causing the spectra to be different. The ratio of the "sample spectrum"
to the "background spectrum" is directly related to the sample's
absorption spectrum.
Accordingly, the technique of "Fourier-transform spectroscopy"
can be used both for measuring emission spectra (for example, the
emission spectrum of a star), and absorption spectra (for example, the absorption spectrum of a liquid).
Continuous-wave Michelson or Fourier-transform spectrograph
The
Fourier-transform spectrometer is just a Michelson interferometer, but
one of the two fully reflecting mirrors is movable, allowing a variable
delay (in the travel time of the light) to be included in one of the
beams.
The Michelson spectrograph is similar to the instrument used in the Michelson–Morley experiment.
Light from the source is split into two beams by a half-silvered
mirror, one is reflected off a fixed mirror and one off a movable
mirror, which introduces a time delay—the Fourier-transform spectrometer
is just a Michelson interferometer with a movable mirror. The beams interfere, allowing the temporal coherence
of the light to be measured at each different time delay setting,
effectively converting the time domain into a spatial coordinate. By
making measurements of the signal at many discrete positions of the
movable mirror, the spectrum can be reconstructed using a Fourier
transform of the temporal coherence
of the light. Michelson spectrographs are capable of very high spectral
resolution observations of very bright sources.
The Michelson or Fourier-transform spectrograph was popular for
infra-red applications at a time when infra-red astronomy only had
single-pixel detectors. Imaging Michelson spectrometers are a
possibility, but in general have been supplanted by imaging Fabry–Pérot instruments, which are easier to construct.
Extracting the spectrum
The intensity as a function of the path length difference (also denoted as retardation) in the interferometer and wavenumber is
where is the spectrum to be determined. Note that it is not necessary for to be modulated by the sample before the interferometer. In fact, most FTIR spectrometers place the sample after the interferometer in the optical path. The total intensity at the detector is
This is just a Fourier cosine transform. The inverse gives us our desired result in terms of the measured quantity :
Pulsed Fourier-transform spectrometer
A pulsed Fourier-transform spectrometer does not employ transmittance techniques.
In the most general description of pulsed FT spectrometry, a sample is
exposed to an energizing event which causes a periodic response. The
frequency of the periodic response, as governed by the field conditions
in the spectrometer, is indicative of the measured properties of the
analyte.
Examples of pulsed Fourier-transform spectrometry
In magnetic spectroscopy (EPR, NMR),
a microwave pulse (EPR) or a radio frequency pulse (NMR) in a strong
ambient magnetic field is used as the energizing event. This turns the
magnetic particles at an angle to the ambient field, resulting in
gyration. The gyrating spins then induce a periodic current in a
detector coil. Each spin exhibits a characteristic frequency of
gyration (relative to the field strength) which reveals information
about the analyte.
In Fourier-transform mass spectrometry,
the energizing event is the injection of the charged sample into the
strong electromagnetic field of a cyclotron. These particles travel in
circles, inducing a current in a fixed coil on one point in their
circle. Each traveling particle exhibits a characteristic cyclotron
frequency-field ratio revealing the masses in the sample.
Free induction decay
Pulsed
FT spectrometry gives the advantage of requiring a single,
time-dependent measurement which can easily deconvolute a set of similar
but distinct signals. The resulting composite signal, is called a free induction decay,
because typically the signal will decay due to inhomogeneities in
sample frequency, or simply unrecoverable loss of signal due to entropic
loss of the property being measured.
Nanoscale spectroscopy with pulsed sources
Pulsed sources allow for the utilization of Fourier-transform spectroscopy principles in scanning near-field optical microscopy techniques. Particularly in nano-FTIR,
where the scattering from a sharp probe-tip is used to perform
spectroscopy of samples with nanoscale spatial resolution, a high-power
illumination from pulsed infrared lasers makes up for a relatively small
scattering efficiency (often < 1%) of the probe.
Stationary forms of Fourier-transform spectrometers
In addition to the scanning forms of Fourier-transform spectrometers, there are a number of stationary or self-scanned forms.
While the analysis of the interferometric output is similar to that of
the typical scanning interferometer, significant differences apply, as
shown in the published analyses. Some stationary forms retain the
Fellgett multiplex advantage, and their use in the spectral region where
detector noise limits apply is similar to the scanning forms of the
FTS. In the photon-noise limited region, the application of stationary
interferometers is dictated by specific consideration for the spectral
region and the application.
One of the most important advantages of Fourier-transform
spectroscopy was shown by P. B. Fellgett, an early advocate of the
method. The Fellgett advantage, also known as the multiplex principle,
states that when obtaining a spectrum when measurement noise is
dominated by detector noise (which is independent of the power of
radiation incident on the detector), a multiplex spectrometer such as a
Fourier-transform spectrometer will produce a relative improvement in
signal-to-noise ratio, compared to an equivalent scanning monochromator, of the order of the square root of m, where m is the number of sample points comprising the spectrum. However, if the detector is shot-noise
dominated, the noise will be proportional to the square root of the
power, thus for a broad boxcar spectrum (continuous broadband source),
the noise is proportional to the square root of m, thus precisely
offset the Fellgett's advantage. For line emission sources the
situation is even worse and there is a distinct `multiplex disadvantage'
as the shot noise from a strong emission component will overwhelm the
fainter components of the spectrum. Shot noise is the main reason
Fourier-transform spectrometry was never popular for ultraviolet (UV)
and visible spectra.
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