Download The History of Infrared Spectroscopy

The History of Infrared Spectroscopy
Robert S. French, HET607, Swinburne Astronomy Online
1. INTRODUCTION
Astronomy has been called the only non-experimental science. The scales of distance, size, time, and energy
are so vast that, except in the simplest of cases, conditions cannot be reproduced in a terrestrial laboratory. Also,
with the exception of planets within our own solar system, we can not send probes to take direct measurements of
astronomical bodies. This makes remote observation all the more important and the invention of the telescope
was undoubtedly the most important development in the history of astronomy. However, spectroscopy, the study
of the wavelengths present in emitted or reflected light, is not far behind.
In 1835, August Compte, a prominent French philosopher, wrote “While we can conceive of the possibility
of determining their [stars’] shapes, their sizes, and their motions, we shall never be able by any means to study
their chemical composition or their mineralogical structure” (Comte 1864). He could not have been more wrong.
It was only 40 years later that spectroscopy was used to discover more than 30 elements in the solar photosphere
(Lockyer 1878). Today, spectroscopy is used to find the composition and temperature of stars, planets, and
interstellar gas, to find extra-solar planets, and to measure the rotation of galaxies, among countless other uses.
While visual-wavelength spectroscopy has the longest history, spectroscopy at other wavelengths from radio to Xray is equally important. In particular, spectroscopy in the infrared (~0.7 m to ~300 m) wavelengths provides
information about the temperature of objects and the composition of complex molecules. In this paper, we will
discuss the development of infrared spectroscopy and the many major discoveries that it has enabled.
2. THE ORIGINS OF INFRARED SPECTROSCOPY
2.1. Early visual spectroscopy
The development of spectroscopy is inextricably linked with advances in the theory of light, refraction, and
diffraction that began in the mid-17th century. Before this time, it was believed that a prism added color to
incoming light. However, Sir Isaac Newton conducted a famous experiment in 1666 in which he used sunlight
and a pair of glass prisms to show that the colors could be dispersed (spread) and then recombined to make white
light (Newton 1672), demonstrating that the colors were already present in the white light. Unfortunately, when a
wide beam of light is shown through a prism, the resulting colors overlap and are smeared and Newton did not
detect any detailed structure in the solar spectrum (Fraunhofer & Ames 1898). It would take nearly 150 years
before the first spectral line was observed in 1802 by William Hyde Wollaston. Wollaston improved upon
Newton’s experiment by using a narrow slit, which reduced the smearing problem, and observed the results from
1012 feet away to allow the spectrum to spread further for greater accuracy (Wollaston 1802). Wollaston
observed five distinct dark lines and two less distinct dark lines in the solar spectrum but did not understand their
significance, thinking they simply delineated the main color bands.
Joseph von Fraunhofer, a German glassmaker, was known for making the highest quality achromatic lenses
and prisms ever seen (Jackson 2000). In 1814, Fraunhofer, unaware of the previous work of Wollaston, observed
the spectra of various flames using a slit and a prism and observed the results with a telescope placed 24 feet away
from the prism, thus making the first spectroscope (Fraunhofer 1817). He found that these spectra contained a
variety of bright lines and that all of them contained the same bright line in the yellow region (which was
identified later as being caused by sodium (Swan 1860)). He continued his experiment by looking at sunlight in
the same manner. Much to his surprise, he discovered an “almost countless” number of dark lines. Trying a
variety of different prisms, Fraunhofer quickly discovered that it was possible to identify particular lines
regardless of the type of glass being used and concluded that the lines were a property of sunlight itself. As a
further experiment, Fraunhofer looked at Venus and observed the same lines as were present in sunlight.
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However, when he looked at several bright stars, including Sirius, he discovered that the lines were very different
from those in sunlight. This was the first use of a spectroscope for astronomical purposes.
2.2. Early infrared spectroscopy
The development of spectroscopy in non-visual wavelengths proceeded in parallel with the development of
visual spectroscopy. Infrared light was discovered by Sir Frederick William Herschel, who performed
experiments with mercury-in-glass thermometers illuminated by sunlight dispersed through a glass prism
(Herschel 1800). Much to his surprise, he found that not only did thermometers register heat beyond the red end
of the visible spectrum, but the greatest amount of heat was found in this region. After an initial claim that these
“heat rays” were just extensions of visible light, he became convinced that they were two separate phenomena.
Thomas Young disagreed. Young had already shown that the “actinic rays” off the blue end of the spectrum
(now called ultraviolet radiation), which had been discovered by Ritter (1803), were as susceptible to diffraction
as visible light and thus represented a continuum of wavelengths. He believed that the “heat rays” would have
similar properties on the other side of the visible. His interpretation was criticized and it was not until nearly 100
years later, when it was shown that visible and infrared light had identical responses to polarization, that the
matter was finally put to rest (Brand 1995).
William Herschel’s son, Sir John Frederick William Herschel, recorded the first infrared spectrum some 40
years after his father’s discovery. He let solar radiation pass through a prism and shine on an alcohol-wetted piece
of paper covered with soot on its back. The alcohol caused the paper to be transparent, allowing the soot to show
through, but dried unevenly, permitting rough observation of the solar spectrum with the alternating white and
black regions on the paper (Herschel 1840). The spectral features were eventually ascribed to the presence of
water vapor in the atmosphere and not to features in the solar spectrum (Brand 1995).
2.3. Dispersion technologies: the prism and the diffraction grating
An important part of an infrared spectrograph is the method used to disperse the incoming light into the
spectrum. The original means of doing so was the prism. A prism works on the principle that the coefficient of
refraction of a material, n, is dependent on the wavelength of the light being refracted, . This causes shorter
wavelengths to bend more at the airprism interface than longer wavelengths, resulting in spectral dispersion.
The angular change (d) of the resulting spectrum with change in wavelength (d) is (Hanel et al. 2003):
d  L dn

d W d
where L is the length of the base of the prism and W is the width of the incoming light beam (Figure 1). The
greater the change of refractive index with wavelength (dn/d), the greater the resolution of the spectrograph.
The resolution, R, is expressed in terms of the wavelength as (Hanel et al. 2003):
R

 L dn

  W d 
where  is the minimum resolvable wavelength interval. Thus, to increase the resolution of a spectrograph, one
can increase the size of the prism, decrease the size of the incoming light beam, or use a material with a greater
change in refractive index with wavelength.
The earliest prisms were made of various types of glass, each having its own refractive properties. Some
types of glass are able to pass wavelengths in the near (~0.75 m) infrared but unfortunately glass, in general, is
opaque to the middle (~540 m) and far (~40300 m) infrared wavelengths. This requires prisms to be made
of other materials. For more than 130 years, rock salt (NaCl) has been the material of choice, with various other
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Figure 1: A prism spectrometer. Paths are shown for two
wavelengths of light,  and . The central and outside rays
are shown for wavelength . Only the central ray is shown for
wavelength  (Hanel et al. 2003).
Figure 2: Illustration of the Huygens
principle and how it causes diffraction
(EIUweb).
salts being used for special applications. However, rock salt is difficult to work with and dissolves easily in
water. Brashear (1886) is credited with being the first person to develop a viable method for creating opticalquality prisms out of rock salt.
A spectrograph must also be properly calibrated so that the location of particular wavelengths in the resulting
spectrum is known. In order for a prism to be used for this purpose, the way in which its index of refraction
changes with wavelength must be carefully measured. Many scientists participated in measuring the refractive
properties of rock salt, including John Herschel, Samuel Langley, William Abney, and William Coblentz (Brand
1995). However, a preferable method is to do away with the prism entirely and to use a different, much more
precise dispersion technique: diffraction.
Diffraction, the bending of waves around corners, can best be understood as an extension of the Huygens
(1690) principle, which considers each point of a wavefront to be the source of a secondary spherical disturbance.
A wavefront thus consists of an infinite number of disturbances, each propagating new disturbances and
advancing the wavefront (Figure 2). When a wavefront encounters a narrow opening, the propagation on the far
side is spherical, not a continuation of the linear wavefront propagation. When the spherical disturbances from
two nearby openings cross, they interfere, alternately reinforcing and canceling each other (Figure 3). This results
in a series of bright and dark bands when projected on a surface. A diffraction grating is based on this principle
but contains many precisely-spaced slits. The large number of slits dramatically increases the sharpness of the
resulting interference maxima, providing a higher resolution spectrum. Diffraction gratings can also be made by
inscribing grooves in a reflective surface, in which case they are called reflection gratings. The physical
mechanism is the same, with the diffraction occurring upon reflection.
Given the angle of incidence, i, the angle of diffraction, r, and the distance between adjacent grooves or slits,
d, light of a particular wavelength, , will be diffracted according to:
m  d (sin i  sin r )
where m is an integer called the order of the diffraction. Because the diffraction is wavelength-dependent,
different wavelengths will be diffracted through different angles, causing the spectrum to be dispersed. Each
order is a copy of the same spectrum, with higher orders being more highly dispersed but lower in intensity
(Figure 4). The orders can overlap one another causing ambiguity in the measured spectrum.
Diffraction had been known to occur for light since the mid-17th century, when Francesco Maria Grimaldi,
an Italian Jesuit priest, studied the phenomenon (Grimaldi 1665; Sommerfeld & Nagem 2004). Newton also
studied diffraction, but explained it using his particle (or “corpuscular”) theory of light. It was Thomas Young
who first offered a viable explanation in his famous double-slit experiment (Young 1804). Young showed that
light passing through two closely spaced, narrow slits produced a series of alternating light and dark bands
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Figure 3: The diffraction caused by Young’s two-slit
experiment. Light travels through slits A and B and is
diffracted. When projected onto a surface, constructive
and destructive interference is visible at points C, D, E,
and F (Diffractionweb).
Figure 4: Illustration of the orders of diffraction for m =
0, 1, and 2. Note how the red-blue angular dispersion is
greater for higher orders, resulting in greater spectral
resolution (HyperPhysicsweb).
(Figure 3). He attributed this to the interference of two waves, an interpretation that was met with incredulity by
the scientific establishment.
The effect that would lead to the invention of the multi-slit diffraction grating was first observed by Robert
Boyle, who noticed in the 17th century that scratches on a glass plate resulted in color in the reflected light.
Young interpreted this as an effect of interference and manufactured the first diffraction grating in 1803 by
carefully inscribing scratches 1/500 inch apart on a glass surface. He measured the resulting fringes and found a
regular mathematical law, but his lack of mathematical sophistication prevented him from codifying it (Brand
1995).
Work on diffraction gratings was continued by Fraunhofer (1821), who was interested in measuring dn/d
for glass so that he could develop more precise optics. Because the dispersion of a diffraction grating depends
only on the slit spacing and not on the properties of any material, diffraction gratings are ideal for precise
wavelength measurement. Fraunhofer manufactured diffraction gratings with slits made from a series of fine
wires on a frame, which lead to the descriptive term “grating”. His experimental skill was impressive, measuring
the distance between wires with a precision of 2 m and the dispersion angles within 4 seconds of arc.
Fraunhofer soon used his gratings to analyze the solar spectrum, and the wavelength measurements he made were
generally within about one part per thousand of modern values. Fraunhofer later switched to engraved glass
surfaces, using diamonds to inscribe the precise grooves. His best known grating sported 302 lines per mm.
Because of Fraunhofer’s extensive work on diffraction gratings, he is usually credited with their invention (Brand
1995).
In addition to the prism and the diffraction grating, it is worth noting one more, extremely simple dispersion
technique: the filter. Filters can be constructed that transmit a well-defined set of wavelengths. If the same
observation is made with multiple filters, the result is essentially a very-low-resolution spectrogram with the
wavelengths measured limited to those transmitted by the filters. While not appropriate for observing many
detailed spectral lines, this technique is nevertheless extremely useful for measuring temperature or observing a
small number of well-defined lines. It also has the advantage of being simple and cheap to implement.
2.4. Early infrared detectors: the thermopile and the bolometer
A photograph of a spectrum can be made by allowing the dispersed light to fall on a photographic plate.
Unfortunately, this technique, which was popular in the visible wavelengths, was poorly suited to the infrared due
to the lack of good infrared-sensitive photochemicals (infrared sensitization techniques would be developed in the
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1920s). As a result, the required preparations were
tedious, the exposures many minutes long, and the
sensitivity limited to the near-infrared wavelengths
(less than 2 m). William Abney was one of the
only scientists able to make productive use of
infrared photography and his technique was so
arcane that no one followed suit (Brand 1995). A
better solution was required.
A single infrared detector, if sufficiently small,
can be placed at one end of a dispersed spectrum
and slowly moved to the other end. The change in
detected intensity with position is a record of the
spectrum. Such a technique was first used with
thermometers, but electrical detectors quickly
became dominant. Thomas Seebeck discovered, in
1822, that connecting two strips of different metals
in a loop and applying a temperature difference to
the two junctions causes a nearby compass to
deflect. He called this the thermomagnetic effect,
which was eventually renamed the thermoelectric
effect once the Danish physicist Hans Christian
Ørsted realized that induced current flow was
Figure 5: Blackbody spectra for a variety of temperatures
(IPSweb).
responsible (Hearnshaw 1996). While a single pair
of metal strips produces only a very small voltage,
multiple junctions can be connected in series to boost the voltage output. Such a device is called a thermopile.
The first practical thermopile was invented by Leopoldo Nobili and Macedonio Melloni. It consisted of 38
bismuth-antimony junctions in a block with a 1 cm2 cross-section. Both sensitivity and reaction time were greatly
improved over thermometers, the previous state of the art. Melloni claimed that he could detect the heat from a
human body at 8 meters. However, the large cross-section of the thermopile limited it to low-resolution
measurements, preventing its use for precision spectroscopy (Brand 1995).
In 1880, Samuel Langley and Frank W. Very took on the task of designing an infrared detector specifically
for dispersive spectroscopy. The result was the bolometer, which consisted of two thin platinum strips covered
with lampblack (Langley 1881). One of the strips was exposed to the incoming radiation and the other was
shielded. The resulting difference in temperature resulted in a small difference in electrical resistance that could
be measured by a galvanometer. Over 20 years, Langley was able to improve the sensitivity of his bolometer by
400 times and he was eventually able to detect the heat from a cow at a quarter mile (Rogalski 2002). Bolometers
are still popular today, especially for detecting the far infrared wavelengths.
2.5. Infrared spectra: experimental and theoretical basis
Spectra come in two flavors: continuum and line. Continuum spectra cover a wide range of wavelengths and
are generally the result of non-quantum processes such as thermal excitation. The most famous continuum
radiation is the Planck blackbody curve. Any object that is above absolute zero will emit radiation (Figure 5) and
the peak wavelength depends on the temperature, with temperatures below ~2,500 K having their peak radiation
in the infrared. Thus infrared spectroscopy is particularly useful for measuring the temperature of cool objects
such as planetary atmospheres or surfaces.
On the other hand, line spectra have sharp, well-defined wavelengths and are the result of quantum processes
either within an atom or between atoms. When Wollaston and Fraunhofer first observed bright lines in the spectra
of flames and dark lines in the solar spectrum they did not know what they meant. It was noted by John Herschel
in 1822 that the bright spectral lines in flames corresponded to the chemicals present in the flame and allowed the
detection of minute quantities of these materials. This was a critical observation that began the science of
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spectroscopy (Thomas 1991). The relationship between the bright and dark lines was finally discovered by
Kirchhoff (1860), who realized that each chemical element had its own distinct set of bright (emission) lines, but
that the same element would produce dark (absorption) lines if a brighter light source was placed behind it.
In 1885, Johann Jakob Balmer examined the spectra of hydrogen and discovered that it emitted a series of
distinct wavelengths. This showed that atoms did not produce a continuum spectrum, although the reason was not
yet understood. Balmer (1885) found that the wavelengths could be described with the formula:

m2 
2
2 
 m n 
  h
where n = 2, h = 3.6456107 m, and integer m > n. Investigations in the infrared led Paschen (1908) to discover
the emissions of hydrogen with n = 3 and Brackett (1922) to discover the emissions with n = 4. Rydberg (1890)
found a similar, more general relationship for ions of other elements that contained only a single electron:
 1
1 
 RZ 2  2  2 

 n1 n2 
1
where R is the Rydberg constant, Z is the atomic number of the element, and integers n1 < n2. These relationships
were finally explained by Bohr (1913), who used the recently proposed quantum theory of Max Planck and Albert
Einstein to conclude that electrons orbited the nuclei of atoms only at defined, or quantized, energy levels. The
electrons can absorb light of a particular wavelength and jump to a higher-energy state, or they can emit a
particular wavelength and return to a lower-energy state. With a few exceptions (such as the Paschen and
Brackett lines of hydrogen), these wavelengths are in the visible portion of the spectrum. Infrared spectral lines
are generally generated by more complex phenomena.
The analysis of infrared spectra began in earnest when Abney & Festing (1882) recorded the spectra of over
50 liquid compounds and found correlations between the absorption lines and the organic groups present in the
molecules. Shortly thereafter, Coblentz (1905) recorded the spectra of hundreds of compounds in the infrared.
The theoretical basis of infrared spectral lines was proposed by Drude (1904), who stated that these lines were due
to vibrations and rotations of molecules caused by their non-zero temperature, not by the oscillations of electrons
within the atoms. At the time, it was believed that the vibration frequencies and rotation velocities would have a
continuous, Maxwellian distribution, resulting in a continuum of spectral features with a single maximum.
However, careful observations by Heinrich Rubens and Eva von Bahr (Imes 1919) showed a series of many
maxima rather than a continuum. It was soon realized that quantum theory had to apply to molecular vibrations
and rotations as well. However, even though the theory of vibrational and rotational spectra is now well
understood, the computations are so complicated that laboratory measurements are still used in most
circumstances to produce reference spectra.
3. RECENT ADVANCES IN INFRARED SPECTROSCOPY
3.1. Semiconductor photon detectors and bolometers
Semiconductors are elements that have a conductivity between that of insulators and metals. They are
usually found in a crystalline lattice, which is integral to their electrical properties. While free electrons are
permitted to have any energy, electrons present within a lattice are restricted in the energies they can possess
much in the same way that electrons in an atom have discrete energy levels. In a semiconductor lattice, the
proximity of other atoms broadens these allowed energy levels into bands. Two of these bands are important
here: the valence band, in which electrons are tightly bound to their host atoms, and the conduction band, in
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which electrons are able to move from atom to atom thus permitting the flow of electric current. Between these is
the forbidden band, otherwise known as the band gap. The conductivity of a semiconductor and the size of the
band gap can by modified by the introduction of trace impurities called dopants. For example, the addition of
only 0.001% of arsenic to a germanium crystal (indicated by Ge:As) will increase the conductivity by a factor of
10,000 (Kruse et al. 1962).
Semiconductors can be doped with elements that provide excess electrons (donors) or elements that bind
excess electrons (acceptors). A semiconductor doped with a donor element is called n-type and one that is doped
with an acceptor element is called p-type. The interface between two adjacent semiconductors, one p-type and
one n-type, forms a potential barrier. Applying an electric field in one direction causes a large current flow, while
applying a field in the opposite direction results in a very small current flow. This p-n junction is the basis of the
diode and, in a more complicated implementation, the basis of the transistor.
Semiconductors are used to detect infrared photons in three main ways. If the interaction of the photons with
the electrons causes them to be raised from the valence band to the conduction band, and this results in a decrease
in the resistivity of the semiconductor, it is called the photoconductive effect. If the interaction of photons with
the electrons at a p-n junction causes a voltage potential to form, it is called the photovoltaic effect. Finally, if the
interaction of photons with the electrons at a p-n junction allows current to flow against the normal direction of
the diode, it is called a photodiode (Kruse et al. 1962). All types of semiconductor photodetectors have fast
response times and excellent signal-to-noise characteristics. However, they are sensitive to thermal excitation and
must be cryogenically cooled for best performance. Cryogenic cooling is often accomplished with superfluid
liquid helium, which achieves temperatures below 3 K (Rogalski 2002). However, cryogenic cooling also
imposes significant disadvantages in weight, size, and cost, and limits mission lifetimes for space-based detectors.
The first infrared photoconductor was developed by Case (1917), who studied the electrical response of a
large number of compounds to light. However, the first practical photoconductor, lead sulfide (PbS), was
discovered by Kutzscher at the University of Berlin. The semiconductor properties of PbS were investigated by
Putley & Arthur (1951), who found a band gap of 1.2 eV. This is equivalent to the energy of a photon with a
wavelength of 1.0 m, which is in the near infrared. The photoconductivity was later found to extend from this
wavelength to 3.2 m (Avery et al. 1954). The next semiconductor to be investigated was indium antimonide
(InSb), with photosensitivity extending to 7.5 m at room temperature. Its sensitivity was one to two orders of
magnitude less than that of a good thermopile, but its response time was 4 million times faster (Avery et al. 1957).
For the past 50 years, HgCdTe has been the preferred material for middle wavelength (330 m) detectors
(Rogalski 2002). HgCdTe has a band gap that is adjustable from 0.7 to 25 m depending on the composition ratio
“x” in Hg1-xCdxTe and the detector temperature (Norton 2002). HgCdTe detectors have a large optical absorption
above the band gap, allowing detectors to be very thin (1020 m), and they can be used in all three ways: as a
photoconductor, a photovoltaic source, or a photodiode. Unfortunately, HgCdTe is a difficult material to
manufacture and work with and has poor structural properties. However, no satisfactory replacement material has
been found to date (Rogalski 2002).
Many other semiconductors have been investigated as infrared detectors. Of particular importance is doped
germanium (Ge), which is sensitive to long wavelengths in the mid-to-far infrared. Germanium detectors can also
be mechanically stressed to lower their band gap energy and thus increase the longest wavelength detectable. For
the longest wavelengths in the far infrared, however, bolometers are still used. Many modern bolometers are
micromachined from a combination of silicon and vanadium dioxide (VO2) and exhibit a change in resistivity of
nearly 4% per Kelvin (Rogalski 2002). Bolometers can also be constructed from pyroelectric materials, which
produce a current flow from spontaneous polarization in the presence of temperature differentials. Pyroelectric
bolometers have been used for infrared spectroscopy on the Pioneer Venus orbiter, the Galileo spacecraft, and the
Mars Exploration Rovers (Lang 2005).
3.2. Michelson interferometers
A diffraction grating works because there is a path difference between two or more light rays from the same
source and they are allowed to constructively or destructively interfere. The amount and type of interference
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depends on the ratio of the path difference to the
wavelength. There is, however, another mechanism for
creating a path difference: the interferometer. An
interferometer takes an incoming beam of light and splits it
into two or more paths with different lengths. The beams
are then recombined and the amount of interference is noted.
The path lengths can be precisely adjusted to measure the
wavelength.
Michelson (1891) was the first scientist to make
extensive use of an interferometer. His device (Figure 6)
Figure 6: The design of a Michelson interferometer
consisted of a half-silvered mirror that split an incoming
(Michelsonweb).
light beam into two paths. The new beams reflected off of
mirrors back to the half-silvered mirror, where they were
recombined and directed towards a detector. One of the
mirrors could be precisely moved to change the
relative path lengths. Michelson used this device, which
now bears his name, to make precise measurements of the
meter using the wavelength of light and to look for the
aether in a famous serious of experiments that laid the
groundwork for Einstein’s special theory of relativity
(Loewenstein 1966).
In the simple case of a single frequency of light (e.g. an
isolated narrow line emission), the frequency of the light can
be determined by simply measuring how far the moveable
mirror has to be moved from one interference peak to the
next. However, when multiple lines or more complex
spectra are present, the dependence of the interference
pattern on the path length difference becomes very
complicated (Figure 7). The most accurate way to derive the
original frequencies from the interferogram is to perform a
Fourier transform. However, computers capable of doing
this processing were not available in the late 19th century
Figure 7: Example of Fourier-transform
and Michelson used a mechanical system of springs and
spectroscopy with ten spectral lines. The
levers to superimpose up to 80 harmonic motions to produce
superimposed light waves, y1  y10, are shown
sample interferograms. The frequencies could be modified
individually. The bottom graph is the linear
until a match with a measured spectrum was found.
combination of the ten waves and illustrates the
The first use of an interferometer to measure infrared
magnitude of interference that would be seen over
radiation was by Rubens & Wood (1911), who used quartz
a range of path length differences in a Michelson
plates as mirrors and recorded the interferogram of the farinterferometer (Hanel et al. 2003).
infrared spectrum of a Welsbach (gas) mantle (found in
modern camping lanterns). They had to guess at the spectral components and they synthesized sample spectra
which could then be matched with the recorded interferogram. Practical use of the interferometer for
spectroscopy would have to wait for the invention of the digital computer.
The first discussion of numerically computed Fourier-transform spectroscopy (FTS) was presented by P. B.
Fellgett in 1951 (Loewenstein 1966). Fellgett also noted that FTS provided a multiplex advantage over standard
dispersion spectroscopy by increasing the signal-to-noise ratio by N , where N is the number of spectral
wavelengths being sampled. This occurs because with a prism or diffraction grating spectrometer with a single
detector, most of the energy in the incoming light is ignored at any given time. However, with FTS all of the
energy from the light is used at all times (Strong & Vanasse 1959).
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The first digitally computed infrared spectrum was produced by H. A. Gebbie, G. A. Vanasse, and J. Strong
in 1956 (Loewenstein 1966) and the first astronomical observation was of the Sun in the far-infrared (Gebbie
1957). Today, with the availability of sensitive infrared detectors and fast digital computers, FTS is a common,
and often preferred, method of spectroscopy.
4. A SAMPLE OF MAJOR INFRARED OBSERVATORIES
4.1. Ground-based telescopes
Dozens of infrared telescopes have been built over the past half century. Due to the strong absorption of
infrared wavelengths by water in Earth’s atmosphere, it is generally desirable to build these telescopes as high as
possible to reduce attenuation, with Mauna Kea in Hawaii, the high plains of Chile, and the American southwest
being ideal locations. Of the major infrared telescopes, the pair of Keck telescopes on Mauna Kea in Hawaii is
undoubtedly among the most well-known. Each Keck telescope consists of a 10 m segmented mirror and a
sophisticated adaptive optics system using both natural and laser guide stars (Keckweb). A variety of cameras
can be attached to each telescope for observation in the visual and infrared wavelengths. As an example,
NIRSPEC (McLean et al. 1998) is a high-resolution (R = 25,000), cryogenically cooled spectrograph that uses a
256256 HgCdTe detector array to cover the near infrared (0.955.1 m).
Although large, dedicated infrared telescopes are useful for detailed analysis of celestial objects, survey
telescopes, which record images and spectra from large sections of the celestial sphere, are best suited for
discovering the objects in the first place. Perhaps the most famous of the celestial surveys, the Sloan Digital Sky
Survey (SDSS) was commissioned in 2000 (York et al. 2000). It uses a dedicated 2.5 m telescope at the Apache
Point Observatory in Sunspot, New Mexico to create a comprehensive survey of 11,600 square degrees of the sky
in five optical bands, two of which are in the near infrared. In addition to visual images taken with a 120megapixel CCD camera (Gunn et al. 1998), it is able to take 640 simultaneous spectra with two digital
spectrographs (Newman et al. 2004). As of the end of 2008, SDSS had cataloged 357 million unique objects and
recorded spectra for over one million galaxies and quasars and 450,000 stars (SDSSweb). SDSS is now in its
third data-gathering phase, which will continue through 2014.
Another extensive survey was the Two Micron All Sky Survey (2MASS), which observed 99.998% of the
celestial sphere in three infrared bands (1.252.16 m) (Skrutskie et al. 2006). Commissioned in 1997, 2MASS
used a pair of 1.3 m telescopes located at Mount Hopkins, Arizona and Cerro Tololo, Chile, to provide the all-sky
coverage. Images were taken using 256256 HgCdTe arrays cooled by liquid nitrogen. 2MASS completed its
mission in 2001 having imaged 471 million objects, of which 1.6 million are galaxies or other extended infrared
sources (Jarrett et al. 2000).
4.2. Space-based telescopes
Although ground-based telescopes have a number of advantages, including lower construction cost, ease of
access, and large mirror sizes, space-based telescopes are nevertheless of particular importance to infrared
observations because of the extensive atmospheric attenuation at even the best locations (Figure 8). Over the past
30 years, a series of ever-more-impressive telescopes have been launched and a comparison of some of their
instruments is shown in Table 1. The first major space-based infrared observatory, the InfraRed Astronomical
Satellite (IRAS) (Neugebauer et al. 1984) was launched into an Earth-based near-polar (99 inclination) orbit on
January 26, 1983. A joint mission of the United States, the Netherlands, and the United Kingdom, IRAS was the
first space-based observatory to perform an all-sky infrared survey. IRAS ran out of helium coolant on November
22, 1983, having completed its mission successfully.
The Infrared Space Observatory (ISO) (Kessler et al. 1996), a mission from the European Space Agency, was
launched into a highly elliptical Earth-based orbit on November 16, 1995. Of particular interest were the Short
-9-
Instrument
Detector Material
Wavelengths (m)
Pixels
Dispersion
Type
1983: Infrared Astronomical Satellite (IRAS)a  0.57 m Aperture
Survey Array
Si:As
16
Filter
8.515
Si:Sb
13
Filter
1930
Ge:Ga
15
Filter
4080
Ge:Ga
15
Filter
83120
Low Resolution Spectrometer
Si:Ga
3
Filter
813
Si:As
2
Filter
1123
1995: Infrared Space Observatory (ISO)  0.60 m Aperture
ISO Camera (ISOCAM)b
InSb
1,024
Filter
2.55.5
Si:Ga
1,024
Filter
418
ISO Short Wavelength Spectrometer (ISO-SWS)c
InSb
48
Grating
2.384.08
Si:Ga
36
Grating
4.0812.0
Si:As
48
Grating
12.029.0
Ge:Be
12
Grating
29.045.2
ISO Long Wavelength Spectrometer (ISO-LWS)d
Ge:Be
1
Grating
4350
Ge:Ga
5
Grating
50110
Stressed Ge:Ga
4
Grating
110190
2003: Spitzer Space Telescope  0.85 m Aperture
Infrared Array Camera (IRAC)e
InSb
Filter
3.195.02
265,536
Si:As
Filter
4.989.34
265,536
Infrared Spectrograph (IRS)f
Si:As
16,384
Grating
5.214.5
Si:As
16,384
Grating
9.919.6
Si:Sb
16,384
Grating
14.038.0
Si:Sb
16,384
Grating
18.737.2
Multiband Infrared Photometer for Spitzer (MIPS)g
Si:As
16,384
Filter
21.526.2
Ge:Ga
1,024
Filter, Grating
62.581.5
Stressed Ge:Ga
40
Filter
139.5174.5
2009: Herschel Space Observatoryh  3.28 m Aperture
Photodetector Array Camera and Spectrometer (PACS)i
Ge:Ga
800
Grating
57210
Spectral and Photometric Imaging REceiver (SPIRE)j
Bolometer (0.3 K)
2
FTS
194671
Table 1: Comparison of four major infrared space telescopes illustrating the increase in spectral range and number
of detectors over time. Sources: a(Neugebauer et al. 1984), b(Cesarsky et al. 1996), c(de Graauw et al. 1996), d(Clegg
et al. 1996), e(Werner et al. 2004), f(IRSHandbookweb), g(MIPSInstHandbookweb), h(Pilbratt et al. 2010),
i
(Poglitsch et al. 2010), j(Griffin et al. 2010).
- 10 -
Wavelength Spectrometer (ISO-SWS) (de
Graauw et al. 1996) and the Long Wavelength
Spectrometer (ISO-LWS) (Clegg et al. 1996).
They each consisted of two scanning grating
spectrometers jointly covering the
wavelengths 2.445 um using arrays of 12
detectors. A full resolution scan took 0.252
hours, depending on the resolution chosen, with
available resolutions running from
R = 1,00025,000.
The Spitzer Space Telescope (Werner et al.
2004), one of NASA’s four Great Observatories,
was launched into an Earth-trailing solar orbit
on August 25, 2003. Three instruments provide
infrared imaging and spectroscopy from 3.6 to
160 m. With low-resolution spectroscopy
instruments, Spitzer was able to achieve
Figure 8: Atmospheric transmission in the infrared on Mauna
Kea, Hawaii, showing the large spectral regions that are
sensitivity comparable with the astrophysical
blocked by Earth’s atmosphere (IACweb).
backgrounds present in the solar system,
particularly emissions from Zodiacal dust.
Among Spitzer’s greatest accomplishments was the first detection of infrared light from an extra-solar planet. In
2009, Spitzer ran out of helium coolant and began the “warm” phase of its mission. The only instrument that can
operate at the warmer (31 K) temperature is IRAC (SpitzerWarmweb).
Herschel, launched May 14, 2009 (Pilbratt et al. 2010), was designed to provide extremely high-resolution
spectroscopy across a wide range of wavelengths and also to extend observations into the far infrared. The
Spectral and Photometric Imaging REceiver (SPIRE) (Griffin et al. 2010) is a Fourier-transform spectrometer
with a bolometer detector cooled to an amazing 0.3 K. Finally, the Wide-field Infrared Survey Explorer (WISE)
(WISEweb) was launched in late 2009. It performed an all-sky survey from 325 m with 500,000 times the
sensitivity of IRAS. While the main mission was completed in July, 2010, and WISE began to run out of solidhydrogen cryogenic coolant shortly thereafter, WISE continues its mission today observing asteroids and comets
within our solar system.
4.3. Interstellar spacecraft
Most interplanetary spacecraft have included infrared spectrometers. Here we will discuss two examples
from spacecraft that have had a major influence on our understanding of the solar system: Voyager and Cassini.
The Voyager 1 and Voyager 2 missions were launched in 1977 (Kohlhase & Penzo 1977). Both identical
spacecraft visited Jupiter and Saturn and Voyager 2 continued on to visit Uranus, and Neptune. Voyager
contained a passively-cooled (200 K) Michelson Fourier-transform infrared spectrometer (IRIS) with a 0.5 m
primary mirror (Hanel et al. 1980). A four-element thermopile allowed detection from 455 m. Raw
interferograms were transmitted to Earth for processing.
The Cassini spacecraft, launched in 1997, entered orbit around Saturn in 2004 and is expected to continue
operating until 2017. It contains two infrared spectrometers. The Composite Infrared Spectrometer (CIRS) is a
Michelson-style interferometer similar to Voyager’s IRIS instrument (Flasar et al. 2004). The use of modern
HgCdTe detectors, however, allows CIRS to operate across a much wider range of wavelengths from the midinfrared to beyond the edge of the infrared spectrum (7 m1000 m). Like IRIS, CIRS transmits raw
interferograms to Earth for processing. The second spectrometer on Cassini, the Visual and Infrared Mapping
Spectrometer (VIMS), operates in the near infrared (visual5.1 m) (Brown et al. 2004). It uses a diffraction
grating along with a scanning mirror to direct spectra onto a 256-element InSb photodetector array. Together,
VIMS and CIRS are able to provide complete coverage of the infrared wavelengths.
- 11 -
5. SOME MAJOR ASTRONOMICAL DISCOVERIES ENABLED BY INFRARED SPECTROSCOPY
5.1. Overview
Infrared spectroscopy, with its ability to measure temperature and detect molecules, provides key
observational insights across a wide range of disciplines in planetary science, astrophysics, and cosmology. With
few exceptions, these observations have been made in the latter half of the 20th century as technology, including
ground-, air-, and space-based telescopes, have become available. In this section we will briefly discuss three of
the many areas in which infrared spectroscopy has made major contributions: the temperature and makeup of
planetary atmospheres, the discovery and analysis of brown dwarfs, and the measurement of the star formation
rate in dusty star-forming galaxies.
5.2. Planetary atmospheres
The temperature of a planet’s atmosphere can be measured with infrared spectroscopy. Comparing the
measured temperature with the equilibrium temperature at the planet’s distance from the Sun can indicate whether
the planet has an internal heat source. In addition, planetary atmospheres are usually composed of molecules,
including organic compounds, making analysis in the infrared especially productive. All planets and satellites
with atmospheres in our solar system have been spectrally analyzed. As a complete review of results for all
planets is beyond the scope of this paper, we will discuss the discoveries regarding the atmosphere of Saturn as an
example.
The temperature of Saturn’s atmosphere was first calculated by Menzel et al. (1926) using the 40-inch
reflector at the Lowell Observatory. They measured the emitted radiation, filtered through water, quartz, glass,
and fluorite (jointly providing wavelengths from 0.3 m to 12.5 m) with a thermocouple and determined a
temperature of 123 K. As the equilibrium temperature caused by solar irradiance is only 65 K, this implied that
Saturn could have an internal heat source (Aumann et al. 1969). Low (1964) measured the brightness temperature
of Saturn with the 82-inch reflector at the McDonald Observatory using wavelengths below 14 m and found a
much lower temperature of 93 K. Aumann et al. (1969) used an airborne observatory flying in the stratosphere at
50,000 feet to measure Saturn’s spectrum from 1.5120 m using a bolometer cooled to 2 K and found a
temperature of 90108 K. They deduced that Saturn must emit four times as much energy as it receives from the
Sun.
More detailed temperature analyses have been provided by the four spacecraft that have been sent to Saturn:
Pioneer 11, Voyager 1 and 2, and Cassini. Orton & Ingersoll (1980) used measurements from Pioneer 11 at 20
and 40 m to determine an equatorial temperature of 96.52.5 K. Most recently, Li et al. (2010) analyzed
Saturn’s temperature profile from both Voyager and Cassini. As previously discovered, they found that Saturn
emits more than twice the energy it receives from the Sun. In addition, they found that during the Voyager era
(approximately one Saturn year ago), Saturn’s northern and southern hemispheres emitted approximately the
same amount of energy, but this is no longer true. Measurements from Cassini showed that Saturn’s southern
hemisphere is now giving off approximately one-sixth more energy than the northern hemisphere. The reason for
this discrepancy is unknown, but may be due to differences in low-level cloud structure.
The infrared spectrum of Saturn was first recorded by Kuiper (1947) using the McDonald Observatory 82inch reflector. Kuiper used a prism and a PbS photodetector to cover the wavelengths 0.752.5 m with R = 80
and identified strong concentrations of methane (CH4) and ammonia (NH3). No further measurements were
reported until Moroz (1962) used the 50-inch reflector at the Crimean Astrophysical Observatory with a
diffraction grating and PbS photodetector. Moroz measured the spectra of Saturn from 0.92.5 m with R = 200
and found significant differences from the spectrum recorded by Kuiper. Because the opening of Saturn’s rings
was at its greatest during the latter observation, Moroz attributed the differences to contamination by the spectra
of the rings, which are composed primarily of water ice.
It would take nearly 20 years before solid evidence for other molecules in Saturn’s atmosphere was reported.
Larson et al. (1980) used a combination of ground- and air-based high-resolution measurements to show that
- 12 -
phosphine (PH3) was present. However, as with measurements of Saturn’s temperature, detailed analysis of
Saturn’s atmospheric composition would best be enabled by spacecraft. Voyager 1 positively identified hydrogen
(H2), helium (He), ammonia, phosphine, methane, ethane (C2H6), and acetylene (C2H2) (Hanel et al. 1981).
Voyager also permitted the detection of deuterated methane (CH3D) and calculation of Saturn’s D/H isotopic ratio
(Courtin et al. 1984). Cassini detected benzene (C6H6), propane (C3H8), and carbon dioxide (CO2). In addition,
Cassini’s observations over several years have allowed more detailed analysis of Saturn’s atmosphere by
comparing observations at different emission angles, which shows different depths in the atmosphere. This has
permitted the calculation of three-dimensional atmospheric models, including the distribution of acetylene and
ethane (Hesman et al. 2009).
5.3. L and T dwarfs
The existence of brown dwarfs, sub-stellar objects that are too low in mass to sustain hydrogen fusion but
nevertheless have fully convective interiors, was first proposed by Kumar (1963). However, it was not until
Becklin & Zuckerman (1988) discovered an excess of infrared radiation around the white dwarf GD 165 that the
existence of brown dwarfs was observationally verified. It took another seven years for the second brown dwarf
to be found, when Nakajima et al. (1995) discovered a companion to the M dwarf Gl 229. Kirkpatrick et al.
(1993) observed GD 165B (the brown dwarf companion to GD 165) with the Low Resolution Imaging
Spectrograph at the Keck Observatory (Keckweb) and found it to have a temperature of ~1900 K and a mass of
~0.07 solar masses, while Saumon et al. (2000) observed Gl 228B with the CGS 4 spectrometer on the United
Kingdom Infrared Telescope (CGS4web) and found it to have a temperature of ~950 K and a mass of
~0.0150.07 solar masses. Gl 228B was sufficiently cool that methane absorption was seen in its atmospheric
spectra. These two objects became the prototypes of two new spectral classes, L and T, respectively, causing the
first significant change to the Harvard stellar classification system in almost 100 years (Leggett et al. 2002).
The search for additional brown dwarfs relied heavily on the infrared survey telescopes and today hundreds
of L dwarfs and tens of T dwarfs are known. More precise spectral classification of L and T dwarfs has been
proposed using near-infrared spectra from 2MASS and SDSS (Burgasser et al. 2002; Leggett et al. 2002; Geballe
et al. 2002; Knapp et al. 2004) and Keck NIRSPEC (McLean et al. 2003). This data was synthesized into a
consistent spectral classification by Cushing et al. (2005), who identified atomic features including neutral Al, Fe,
Mg, Ca, Ti, Na, and K and molecular features including CH4, H2O, and VO that could be used in the classification
process.
As essentially very large gas giant planets, L and T dwarfs have atmospheres and associated atmospheric
dynamics including convection, winds, and clouds. In addition to classification, infrared spectra have been used
to analyze the atmospheres of L and T dwarfs, providing information about the correlation of neutral alkali metals
with atmospheric pressure (Burrows et al. 2000) and the processes of cloud formation, sedimentation, and
atmospheric chemical equilibrium (Marley et al. 2002).
5.4. Dusty star forming galaxies
The mass of galaxies near the Milky Way is concentrated in stars, with less than 10% of the mass present in
the form of gas or dust. These stars formed from dense molecular clouds and it is assumed that the evolution of
other galaxies proceeded in a similar fashion. Thus the measurement of star formation rates (SFR) over the life of
a galaxy would provide important information about galactic evolution. While we can not watch a single galaxy
for that long, we can nevertheless achieve the same goal by observing a variety of redshifted galaxies of different
ages and thus in different stages of their evolution.
Star formation rates are often estimated by looking at ultraviolet (UV) and visual emissions since galaxies
undergoing rapid star formation are dominated by hot O and B class stars, which emit light primarily in these
wavelengths. However, since these stars form within dusty molecular clouds, much of the evidence of their star
formation is hidden and surveys that use only UV and visual emissions can dramatically underestimate the rate of
star formation. Galaxies undergoing bursts of star formation within dusty clouds would be expected to emit
significant amounts of infrared light as the dust is heated by the stars to ~30 K and then reradiated in the infrared
- 13 -
(Hughes et al. 1998). In these galaxies, infrared emission dominates emission at visual wavelengths, sometimes
by as much as a factor of ten. Proper modeling of the SFR can be performed using a combination of ultraviolet,
visual, and infrared observations.
The first emissions from infrared bright galaxies were observed by Kleinmann & Low (1970). Rieke & Low
(1972) expanded upon these observations and presented information about 32 infrared sources. The first
estimates of star formation rates in infrared bright galaxies were made by Thronson & Telesco (1986). They
assumed that the infrared emission was only caused by stars that formed in the very recent past (~2106 years)
because the radiation pressure from these new stars along with supernovae would quickly dissipate the
surrounding dusty molecular cloud. Combining this assumption with a Salpeter initial mass function, they
derived a star formation rate of:
M * / M   6.5 1010 LIR / L
where M * is the star formation rate per year and LIR is the galaxy’s infrared luminosity. Likewise, the long-term
(0.46109 year) star formation rate can be derived from the blue luminosity, LB. The combination of these two
equations permits the comparison of short-term and long-term star formation rates. Thronson & Telesco found
six galaxies using IRAS where the short-term star formation rate exceeded the long-term rate, implying they were
currently undergoing a starburst event. Rowan-Robinson (1997) improved upon this result by analyzing infrared
galaxies observed by ISO in the Hubble Deep Field and found that they were undergoing prodigious star
formation as well.
Once a method of calculating the star formation rate is found, it is then possible to analyze the change in star
formation over time by correlating it with redshift. Blain et al. (1999) used galaxies detected by IRAS to fit the
relative comoving density of star formation to the function (1  z ) p . They found a best-fit dust temperature of
38 K and p = 3.8 for small z. Spitzer observed ~2,600 infrared sources with the MIPS instrument and found that
the comoving energy density of the universe changes with (1  z )3.9 0.4 up to z ~ 1, and the infrared bright
galaxies are responsible for 70%15% of this energy at z ~ 1. Given that the energy contribution from UV
sources is known to evolve as (1  z ) 2.5 , infrared bright galaxies dominate star-forming activity for z > 0.7 (Le
Floc’h et al. 2005).
6. NEW AND FUTURE TECHNOLOGY
Infrared technology continues to improve and infrared observations have become even more important with
the current research emphasis on galactic evolution and brown dwarfs. As a result, many new infrared
observatories and spacecraft are being built. The largest and most impressive of these is the James Web Space
Telescope (JWSTweb). Billed as a replacement for the Hubble Space Telescope (HST), the JWST, with its 6.5 m
mirror, will be the largest telescope ever launched into orbit. As its mirror diameter is 2.7 times that of the HST,
the JWST will be able to collect more than seven times as much light. Two infrared sensors, one covering
0.65 m and the other covering 529 m, provide a spectral range significantly greater than the 1.02.5 m
range of the HST’s NICMOS camera (Thompson 1992). The greater light-gathering capability and sensitivity to
longer wavelengths will together allow the observation of much dimmer and more highly redshifted objects. The
JWST is currently expected to launch in 2014 and will be placed in an elliptical orbit around the Sun-Earth L2
Lagrange point so that it can be simultaneously shielded from infrared radiation emitted from the Sun, Earth, and
Moon.
Another new observing platform is SOFIA, the Stratospheric Observatory for Infrared Astronomy
(SOFIAweb). SOFIA consists of a Boeing 747SP aircraft and a 2.5-meter reflecting telescope that is pointed
through a large hole in the aircraft’s fuselage. The telescope’s location in an aircraft allows for high altitude
(greater than 39,000 feet, above 99% of the water in the atmosphere) observation, all-sky coverage, and easy
- 14 -
maintenance and instrument replacement. SOFIA completed its first astronomical observation on December 1,
2010.
7. CONCLUSION
The principles of infrared spectroscopy, including methods of spectral dispersion (prisms and diffraction
gratings) and detection (thermopiles and bolometers) have been known since the mid-19th century. Even though
the first infrared spectrum of the Sun was recorded in 1840, it has only been since the development of
semiconductor-based detectors and the construction of high altitude, airborne, or space-based observatories in the
second half of the 20th century that infrared spectroscopy has become a significant contributor to astronomical
research. It would be difficult to understate the importance of infrared observations to our understanding of the
solar system, Galaxy, and universe. Infrared spectroscopy has been used to study the temperatures and
composition of planets, to detect and analyze brown dwarfs, and to better understand the history of galactic
evolution and star formation, among many other things. With the current emphasis on learning about star
formation and galactic evolution, the importance of infrared astronomy will only increase over the next decades.
Luckily, there is an impressive lineup of new telescopes that will be used to further this research, and infrared
astronomy will continue to yield exciting results for many years to come.
References
Abney, C., & Festing, L. 1882, Phil. Trans., 172, 887
Aumann, H. H., Gillespie, C. M., & Low, F. J. 1969, BAAS, 213
Avery, D. G., Goodwin, D. W., Lawson, W. D., & Moss, T. S. 1954, Proc. Phys. Soc. B, 67, 761
Avery, D. G., Goodwin, D. W., & Rennie, M. A. E. 1957, J. Sci. Instrum., 34, 394
Balmer, J. J. 1885, Ann. Phys., 261, 80
Becklin, E. E., & Zuckerman, B. 1988, Nature, 336, 656
Blain, A. W., Smail, I., Ivison, R. J., & Kneib, J. 1999, MNRAS, 302, 632
Bohr, N. 1913, Phil. Mag., 26, 1
Brackett, F. S. 1922, ApJ, 56, 154
Brand, J. C. D. 1995, Lines of light: the sources of dispersive spectroscopy, 1800-1930 (Amsterdam: OPA)
Brashear, J. A. 1886, The Sidereal Messenger, 5.5, 149
Brown, R. H. et al. 2004, SSR, 115, 111
Burgasser, A. J. et al. 2002, ApJ, 564, 421
Burrows, A., Marley, M. S., & Sharp, C. M. 2000, ApJ, 531, 438
Case, T. W. 1917, Phys. Rev., 9, 305
Cesarsky, C. J. et al. 1996, A&A, 315, L32
CGS4web: CGS4, http://www.jach.hawaii.edu/UKIRT/instruments/cgs4/cgs4.html (accessed 11 November 2010)
Clegg, P. E. et al. 1996, A&A, 315, L38
Coblentz, W. W. 1905, Investigations of Infra-Red Spectra (Washington, D.C.: Carnegie Institute of Washington)
Comte, A. 1864, Cours de Philosophie Positive (Paris: J. B. Baillière et Fils)
Courtin, R., Gautier, D., Marten, A., Bezard, B., & Hanel, R. 1984, ApJ, 287, 899
Cushing, M. C., Rayner, J. T., & Vacca, W. D. 2005, ApJ, 623, 1115
Diffractionweb: Young Diffraction, http://upload.wikimedia.org/wikipedia/commons/8/8a/Young_Diffraction.png
(accessed 10 November 2010)
Drude, P. 1904, Ann. Phys., 14, 677
EIUweb: Wave Optics, http://www.ux1.eiu.edu/~cfadd/1160/Ch25WO/Huygn.html (accessed 20 November
2010)
Flasar, F. M. et al. 2004, SSR, 115, 169
Fraunhofer, J. 1817, Denkschr. K. Akad. Wiss., München, 5, 193 (translation from Fraunhofer & Ames 1898)
Fraunhofer, J. 1821, Denkschr. K. Akad. Wiss., München, 8, 1 (translation from Fraunhofer & Ames 1898)
- 15 -
Fraunhofer, J. V., & Ames, J. S. 1898, Prismatic and diffraction spectra: Memoirs by Joseph von Fraunhofer
(New York: Harper & Brothers)
Geballe, T. R. et al. 2002, ApJ, 564, 466
Gebbie, H. A. 1957, Phys. Rev., 107, 1194
de Graauw, T. et al. 1996, A&A, 315, L49
Griffin, M. J. et al. 2010, A&A, 518, 7
Grimaldi, F. M. 1665, Physico-mathesis de lvmine, coloribvs, et iride, aliisque adnexis libri duo: opvs posthvmvm
Gunn, J. E. et al. 1998, AJ, 116, 3040
Hanel, R. et al. 1981, Science, 212, 192
Hanel, R., Crosby, D., Herath, L., Vanous, D., Collins, D., Creswick, H., Harris, C., & Rhodes, M. 1980, Appl.
Opt., 19, 1391
Hanel, R. A., Conrath, B. J., Jennings, D. E., & Samuelson, R. E. 2003, Exploration of the Solar System by
Infrared Remote Sensing (2nd ed.; Cambridge: Cambridge University Press)
Hearnshaw, J. B. 1996, The measurement of starlight: two centuries of astronomical photometry (Cambridge:
Cambridge University Press)
Herschel, J. F. W. 1840, Phil. Trans. Royal Soc. London, 130, 1
Herschel, W. 1800, Phil. Trans. Royal Soc. London, 90, 255
Hesman, B. E. et al. 2009, Icarus, 202, 249
Hughes, D. H. et al. 1998, Nature, 394, 241
Huygens, C. 1690, Treatise on Light (Chicago: University of Chicago Press)
HyperPhysicsweb: Diffraction Grating, http://hyperphysics.phy-astr.gsu.edu/hbase/phyopt/grating.html (accessed
20 November 2010)
IACweb: How the atmosphere affects astronomy with CanariCam,
http://www.iac.es/project/CCam/Atmosphere.htm (accessed 10 November 2010)
Imes, E. S. 1919, ApJ, 50, 251
IPSweb: New paints and coatings with specific properties in the near and the far infrared, http://www.ipsinnovations.com/new_paints_ref.htm (accessed 29 November 2010)
IRSHandbookweb: Spitzer: IRS Instrument Handbook, http://ssc.spitzer.caltech.edu/irs/irsinstrumenthandbook/
(accessed 24 November 2010)
Jackson, M. W. 2000, Spectrum of Belief: Joseph von Fraunhofer and the Craft of Precision Optics (Cambridge,
MA: The MIT Press)
Jarrett, T. H., Chester, T., Cutri, R., Schneider, S., Skrutskie, M., & Huchra, J. P. 2000, AJ, 119, 2498
JWSTweb: The James Webb Space Telescope, http://www.jwst.nasa.gov/index.html (accessed 12 November
2010)
Keckweb: W. M. Keck Observatory, http://www.keckobservatory.org/ (accessed 3 December 2010)
Kessler, M. F. et al. 1996, A&A, 315, L27
Kirchhoff, G. 1860, Phil. Mag. Ser. 4, 20(130), 1
Kirkpatrick, J. D., Henry, T. J., & Liebert, J. 1993, ApJ, 406, 701
Kleinmann, D. E., & Low, F. J. 1970, ApJ, 159, L165
Knapp, G. R. et al. 2004, AJ, 127, 3553
Kohlhase, C. E., & Penzo, P. A. 1977, SSR, 21, 77
Kruse, P. W., McGlauchlin, L. D., & McQuistan, R. B. 1962, Elements of Infrared Technology: Generation,
Transmission, and Detection (New York: John Wiley & Sons, Inc.)
Kuiper, G. P. 1947, ApJ, 106, 251
Kumar, S. S. 1963, ApJ, 137, 1121
Lang, S. B. 2005, Phys. Today, 58(8), 31
Langley, S. P. 1881, Proc. Am. Acad. Arts. Sci., 16, 342
Larson, H. P., Fink, U., Smith, H. A., & Davis, D. S. 1980, ApJ, 240, 327
Le Floc’h, E. et al. 2005, ApJ, 632, 169
Leggett, S. K. et al. 2002, ApJ, 564, 452
Li, L. et al. 2010, JGR (Planets), 115, 11002
- 16 -
Lockyer, J. N. 1878, Proc. Royal Soc. London, 28, 157
Loewenstein, E. V. 1966, Appl. Opt., 5, 845
Low, F. J. 1964, ApJ, 69, 550
Marley, M. S., Seager, S., Saumon, D., Lodders, K., Ackerman, A. S., Freedman, R. S., & Fan, X. 2002, ApJ,
568, 335
McLean, I. S. et al. 1998, in Proc. SPIE Vol. 3354, Infrared Astronomical Instrumentation, ed. M. Fowler, 566
McLean, I. S., McGovern, M. R., Burgasser, A. J., Kirkpatrick, J. D., Prato, L., & Kim, S. S. 2003, ApJ, 596, 561
Menzel, D. H., Coblentz, W. W., & Lampland, C. O. 1926, ApJ, 63, 177
Michelson, A. A. 1891, Phil. Mag., 31, 256
Michelsonweb: Fourier Transform Spectrometer,
http://scienceworld.wolfram.com/physics/FourierTransformSpectrometer.html (accessed 23 November
2010)
MIPSInstHandbookweb: Spitzer: MIPS Instrument Handbook,
http://ssc.spitzer.caltech.edu/mips/mipsinstrumenthandbook/ (accessed 24 November 2010)
Moroz, V. I. 1962, Sov. Astron., 5, 827
Nakajima, T., Oppenheimer, B. R., Kulkarni, S. R., Golimowski, D. A., Matthews, K., & Durrance, S. T. 1995,
Nature, 378, 463
Neugebauer, G. et al. 1984, ApJ, 278, L1
Newman, P. R. et al. 2004, preprint (astro-ph/0408167)
Newton, I. 1672, Phil. Trans. Royal Soc. London, 5, 3075
Norton, P. 2002, Opto-Elec. Rev., 10, 159
Orton, G. S., & Ingersoll, A. P. 1980, JGR, 85, 5871
Paschen, F. 1908, Ann. Phys., 332(13), 537
Pilbratt, G. L. et al. 2010, A&A, 518, 6
Poglitsch, A. et al. 2010, A&A, 518, 12
Putley, E. H., & Arthur, J. B. 1951, Proc. Phys. Soc. B, 64, 616
Rieke, G. H., & Low, F. J. 1972, ApJ, 176, L95
Ritter, J. W. 1803, Ann. Phys., 12(12), 409
Rogalski, A. 2002, IR Phys. Tech., 43(3-5), 187
Rowan-Robinson, T. I. C. M. 1997, preprint (astro-ph/9707030)
Rubens, H., & Wood, R. W. 1911, Phil. Mag., 21, 249
Rydberg, J. R. 1890, Phil. Mag., 29, 331
Saumon, D., Geballe, T. R., Leggett, S. K., Marley, M. S., Freedman, R. S., Lodders, K., Fegley, B., & Sengupta,
S. K. 2000, ApJ, 541, 374
SDSSweb: SDSS Data Release 7, http://www.sdss.org/dr7/ (accessed 1 December 2010)
Skrutskie, M. F. et al. 2006, ApJ, 131, 1163
SOFIAweb: SOFIA, http://www.nasa.gov/mission_pages/SOFIA/index.html (accessed 12 November 2010)
Sommerfeld, A., & Nagem, R. J. 2004, Mathematical theory of diffraction (Boston: Birkhäuser)
SpitzerWarmweb: NASA's Spitzer Telescope Warms Up To New Career,
http://www.nasa.gov/mission_pages/spitzer/news/spitzer-20090506.html (accessed 13 November 2010)
Strong, J., & Vanasse, G. A. 1959, J. Opt. Soc. Am., 49, 844
Swan, W. 1860, Phil. Mag. Ser. 4, 20(130), 173
Thomas, N. C. 1991, J. Chem. Ed., 68(8), 631
Thompson, R. 1992, SSR, 61, 69
Thronson, H. A., & Telesco, C. M. 1986, ApJ, 311, 98
Werner, M. W. et al. 2004, ApJS, 154, 1
WISEweb: Wide-Field Infrared Survey Explorer, http://www.nasa.gov/mission_pages/WISE/main/index.html
(accessed 12 November 2010)
Wollaston, W. H. 1802, Phil. Trans. Royal Soc. London, 92, 365
York, D. G. et al. 2000, ApJ, 120, 1579
Young, T. 1804, Phil. Trans. Royal Soc. London, 94, 1
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