In the previous chapter we saw how the ultraviolet (UV), visible, infrared (IR), and microwave spectra of the giant planets are formed, and how the absorption features of different gases (and theoretically aerosols) are visible in these spectra. Clearly much can be learned about the atmospheres from observing the spectra of these planets, and in this chapter we will review the measurements that have been made to date and how they may be used.
In this chapter we will review some of the technical details of measuring the UV-microwave spectra of the planets and in Section 7.2 we will briefly review how such radiation is detected and how spectra are measured. Prior to 1973, the only measurements of the giant planets that were available were telescope observations from the surface of the Earth in the visible, IR, and microwave wavelengths. Such observations have the obvious attraction that they are relatively easy to do, and have a number of other advantages, although there are drawbacks as we shall see in Section 7.3. We shall also see that the technology of detection and data processing has improved dramatically over the years and thus ground-based observations continue to be a rich source of information on the giant planets to this day. In Sections 7.4 and 7.5 we will look at some of the major ground and airborne visible/IR telescope facilities around the world that are engaged in outer planet observations and in Section 7.6 we will look at ground-based microwave observations. Many of the problems encountered by ground-based visible/IR telescopes are negated by placing the telescope in orbit around the Earth, and thus in Section 7.7 we will look at recent space telescopes such as the Hubble Space Telescope (HST), the Infrared Space Observatory (ISO), and the Spitzer Space Telescope.
In 1973, the first spacecraft mission arrived at Jupiter and since then a number of spacecraft have flown past the giant planets, and more recently have been placed in orbit to conduct extended campaigns. These missions have enormous advantages over ground-based and space-based telescopes, but are of course immensely expensive and difficult to achieve. In Section 7.8 we will review the flyby missions of the giant planets and then in Section 7.9 we will consider the orbiting missions of Galileo and Cassini/Huygens. Finally we will discuss retrieval methods in Section 7.10 where we will see how remotely sensed observations are actually used to infer atmospheric properties in planetary atmospheres.
7.2 MEASUREMENT OF VISIBLE, IR, AND MICROWAVE SPECTRA
Before going on to look at the current sources of remotely sensed spectral data on the giant planets and how these data are reduced to infer atmospheric properties, we will briefly look at how the spectra of these planets are actually measured from the visible through to microwave wavelengths.
For all remote-sensing instruments, the incident radiance must first be collected and focused onto detecting elements in order to record a signal. In any detection system there are a number of sources of noise and the design of remote-sensing instruments aims to minimize these in order to achieve a high signal-to-noise ratio. Suppose that radiance B (Wm~2 sr-1(cm-1)-1) is incident upon a remote-sensing system. The power P(W) incident on the detector is given by where F(v) is the spectral transmission of the optical system; and the AO product is given either by the area of the entrance aperture multiplied by the solid angle of the field of view (FOV) observed by the instrument, or equivalently by the area of the detector multiplied by the solid angle of the cone of radiation condensed onto the detector by the instrument optics. The quantity E in Equation (7.1) refers to thermal radiation self-emitted by the telescope and optics, which is incident on the detector. The signal detected by the instrument (e.g., volts or amps) is then given by where R(v, T) is the spectral responsivity of the detector, and has units of V/W or A/ W. In the detection and pre-amplification stages of instruments there are numerous sources of noise such as Johnson noise (or voltage noise, which appears across resistances), Shot noise (or current noise, arising from the fact that a "steady current" is actually composed of a stream of individual electrons), noise arising from the incident radiation itself, and radiation from the optical elements if we are considering the thermal-infrared (e.g., Hanel et al., 2003). The Shot noise associated with a current 70 is given by 1S = 2e/0A/, where A/ is the bandwidth and e is the electron charge, and is clearly minimized by limiting the currents in the detection stages of
amplification. Johnson noise, however, given by V] = 4kBTRAf, where R is the resistance and kB is the Boltzmann constant, depends on temperature, as does the noise of radiation thermally emitted by the optics and filters. Hence, to maximize the signal-to-noise ratio of the detected radiance, especially when working in the thermal-infrared, the detectors, filters, and as much of the telescope optical system as is possible must be cooled to low temperatures. Adding all sources of noise together, a common figure of merit of the instrument is the noise equivalent power (NEP), which is defined as the power of incident radiation that, when viewed in a 1 Hz bandwidth, gives a signal equivalent to all the sources of noise. Another very useful figure of merit for comparing the sensitivities of detector is D * which is defined as
NEP V J
where A is the area of the detector. Highly sensitive detectors have a high D *. Finally, for thermal-IR observations, all the sources of noise may be analyzed in terms of their noise equivalent radiance (NER), defined as the incident spectral radiance which gives a signal-to-noise ratio of unity.
There are two main ways of detecting radiation: (1) photon detectors, which detect individual photons; and (2) thermal detectors, or bolometers, which detect the temperature rise of elements exposed to radiation. Examples of photon detectors are: photovoltaic cells, where absorbed photons promote the production of an electron-hole pair in a p-n junction and thus produce a transient voltage; and photoconductive detectors, where absorbed photons again promote the production of an electron-hole pair in an element of semiconductor material, which temporarily alters its conductivity. Examples of thermal conductors include: thermopiles, which are basically a stacked array of bimetallic junctions that produce a voltage dependent on their temperature via the thermocouple effect; and pyroelectric detectors, which use a dielectric material with a temperature-sensitive dipole moment sandwiched between the plates of a capacitor. Absorption of thermal radiation modifies the permittivity and thus the capacitance. The choice of detector depends on many things: cost, required signal-to-noise ratio, and response time. The reader is referred to more specialized texts for further information (Hanel et al., 2003; Houghton and Smith, 1966; Smith et al., 1968).
Many remote-sensing instruments simply record the incident radiation received within a bandwidth defined by a set of spectral filters, and most imaging cameras operate in this way. Photometers record accurately the flux level of visible or near-IR light within narrow spectral channels, while radiometers perform a similar function at thermal-IR wavelengths. Such an instrument design is cheap, reliable, and ideally suited to imaging although it does require a priori knowledge of the planetary spectrum in order to place the channel filters at suitable wavelengths.
Where the planetary spectrum is less well known, spectrometers must be used of which in the visible/IR there are two main types: grating spectrometers and interferometers. For really high-resolution visible/IR work, Fabry-Perot interferometers may also be used, usually in conjunction with a grating spectrometer, which correctly limits the range of wavelengths that are passed through the Fabry-Perot interferometer to eliminate aliasing.
Grating spectrometers are often used for IR spectroscopy, particularly at near-IR wavelengths and have the advantage of relative simplicity. Light is collected from the planet via a telescope system and then the collimated light illuminates a reflecting diffraction grating, which disperses a spectrum onto the focal plane (Figure 7.1). The grating is "blazed" to maximize the throughput at the central wavelength of the region of interest. In its simplest form, there is a single IR detector at the focal plane, and the spectrum is measured by recording the detector signal as the grating is scanned through a small angle. Since IR detectors are usually sensitive to a wide range of wavelengths, an "order-sorting" filter must in practice also be added to limit the range of wavelengths that can be detected so that the spectrometer only operates in the spectral order desired.
While such a design is simple, only a small range of wavelengths are recorded at a time, and thus much of the radiation that is dispersed by the grating is wasted. Also it takes a certain length of time to scan the grating through the angular range required
to build up a spectrum, and this can be particularly problematic from an observational point of view when the instrument is observing a planet whose spectrum alters significantly with position. During the time the grating is scanned, the spectrometer may be looking at a very different region at the end of the scan to that observed at the beginning, and thus the recorded spectrum may be very hard to interpret. An alternative approach, sometimes called a spectrograph, is to have a whole array of closely spaced, contiguous detectors in the focal plane, each recording different wavelengths, and thus leave the grating angle fixed. While such instruments are much harder and more expensive to build, a spectrum is recorded much more quickly, and if the region of the planet observed should vary, all parts of the recorded spectrum are equally affected. An intermediate approach is to use fewer detectors, spaced farther apart on the focal plane and then scan the grating over a short range, such that the final spectrum is built up from a number of subspectra recorded by each individual detector. This design is easier to build, records a spectrum in a reasonably short space of time, but can suffer from variable scene problems which reveal themselves as mismatches between the individual detector subspectra where they overlap.
Clearly, for planetary work it is of great interest to record the spectrum at a number of locations on the planet in order to build up a multispectral image. This can be achieved by scanning the instrument in both x and y directions, but in more recent years it has been possible to construct imaging spectrographs where a two-dimensional array of detectors is placed in the focal plane, and thus multiple spectra are simultaneously recorded from a small range of viewing angles perpendicular to the grating dispersion direction. The instrument then need only be scanned in one direction (the grating dispersion direction)—not two—in order to construct a multispectral image.
The spectral resolution of a grating spectrometer is fixed by the size of the entrance aperture to the collimator, the dispersion of the grating, and by the physical dimensions of the detectors. For higher resolution spectroscopy, some instrument designs allow for a Fabry-Perot interferometer to enter the beam as indicated in Figure 7.1. This allows very high-resolution spectrometry to be conducted over a small spectral range with the resolving power (A/AA) of the order of 10,000. Other designs of grating spectrometers use Echelle spectroscopy, where the spectrum is dispersed in two directions onto a two-dimensional pixel array. In this way many orders of interference are observed simultaneously, allowing an Echelle spectrometer to have very much higher resolution than a conventional grating spectrometer.
While grating spectrometers perform spectroscopy by "division of wavefront", interferometers are also often used at mid-IR to far-IR wavelengths, which operate by division of amplitude. The simplest example is a Michelson interferometer, where light is split into two beams by a beam-splitting mirror and then recombined onto a detector (as in Figure 7.2). For monochromatic light of intensity 70, the intensity at the detector varies with scanning mirror position x as (e.g., Hecht and
Figure 7.2. Michelson interferometer layout.
Figure 7.2. Michelson interferometer layout.
where k = 2^/A, and A = 2x. For incident light with a spread of wavelengths and spectral density I(k) dk, Equation (7.4) must be modified to
and since IR detectors are typically sensitive to a wide range of wavelengths A£, the detected intensity must be integrated over all wavelengths in this range:
Hence, in the limit that Av is large it can be seen that the intensity detected is equal to a mean value plus the Fourier cosine transform of the incident spectrum. The signal I (A), recorded as a function of mirror position, is called an interferogram and the incident spectrum may then theoretically be perfectly reconstructed by the inverse Fourier cosine transform as
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