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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 1012 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. -1- 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 airprism 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 (~0.75 m) infrared but unfortunately glass, in general, is opaque to the middle (~540 m) and far (~40300 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 -2- 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 -3- 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 -4- 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 -5- 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 = 3.6456107 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 -6- 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 (330 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 (1020 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 -7- 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). -8- 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 256256 HgCdTe detector array to cover the near infrared (0.955.1 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 (1.252.16 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 256256 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 8.515 Si:Sb 13 Filter 1930 Ge:Ga 15 Filter 4080 Ge:Ga 15 Filter 83120 Low Resolution Spectrometer Si:Ga 3 Filter 813 Si:As 2 Filter 1123 1995: Infrared Space Observatory (ISO) 0.60 m Aperture ISO Camera (ISOCAM)b InSb 1,024 Filter 2.55.5 Si:Ga 1,024 Filter 418 ISO Short Wavelength Spectrometer (ISO-SWS)c InSb 48 Grating 2.384.08 Si:Ga 36 Grating 4.0812.0 Si:As 48 Grating 12.029.0 Ge:Be 12 Grating 29.045.2 ISO Long Wavelength Spectrometer (ISO-LWS)d Ge:Be 1 Grating 4350 Ge:Ga 5 Grating 50110 Stressed Ge:Ga 4 Grating 110190 2003: Spitzer Space Telescope 0.85 m Aperture Infrared Array Camera (IRAC)e InSb Filter 3.195.02 265,536 Si:As Filter 4.989.34 265,536 Infrared Spectrograph (IRS)f Si:As 16,384 Grating 5.214.5 Si:As 16,384 Grating 9.919.6 Si:Sb 16,384 Grating 14.038.0 Si:Sb 16,384 Grating 18.737.2 Multiband Infrared Photometer for Spitzer (MIPS)g Si:As 16,384 Filter 21.526.2 Ge:Ga 1,024 Filter, Grating 62.581.5 Stressed Ge:Ga 40 Filter 139.5174.5 2009: Herschel Space Observatoryh 3.28 m Aperture Photodetector Array Camera and Spectrometer (PACS)i Ge:Ga 800 Grating 57210 Spectral and Photometric Imaging REceiver (SPIRE)j Bolometer (0.3 K) 2 FTS 194671 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 2.445 um using arrays of 12 detectors. A full resolution scan took 0.252 hours, depending on the resolution chosen, with available resolutions running from R = 1,00025,000. 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 325 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 455 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 m1000 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 (visual5.1 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 1.5120 m using a bolometer cooled to 2 K and found a temperature of 90108 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 96.52.5 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 0.752.5 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 0.92.5 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 ~0.0150.07 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 (~2106 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 1010 LIR / L where M * is the star formation rate per year and LIR is the galaxy’s infrared luminosity. Likewise, the long-term (0.46109 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 0.65 m and the other covering 529 m, provide a spectral range significantly greater than the 1.02.5 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. 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