Photoconductors and photodiodes are combined with electronic readouts to make detector arrays. The progress in integrated circuit design and fabrication techniques has resulted in continued rapid growth in the size and performance of these solid state arrays. In the infrared, these devices are based on a (conceptually) simple combination of components that have already been discussed; an array of integrating amplifiers is connected to an array of detectors. These devices can provide thousands to millions of detector elements, each of which performs near the fundamental photon noise limit for most applications between wavelengths of 1 and 40 μm. For visible and near-infrared detection, monolithic structures can be built in silicon to form arrays with millions of high performance pixels. The most highly developed of these visible detectors is the charge coupled device, or CCD, which combines an array of intrinsic photoconductors, a sequentially addressable array of integrating capacitors, and a FET output amplifier, all built together on a single wafer of silicon. Operated at low temperatures, CCDs can reach fundamental detection limits for virtually all applications between wavelengths of ∼0.1 nm (in the X-ray region) and ∼1 μm (in the near infrared), except when rapid time response is required. So-called CMOS imaging arrays are an alternative to CCDs.

More and more investigators are being attracted to work in the infrared spectral region. My intention in this book is to introduce the infrared, its peculiarities and its special techniques, to a new audience that may include established astronomers as well as new students. Basic facts and figures are emphasized and I have not tried to cover all current research. The approach usually avoids the historical and the reader is referred instead to recent books and papers. The origins of ideas are always complex and, with something like 104 papers published, it is impossible to do justice to every individual in the field.

The first chapter deals with the basic facts of blackbody radiation, which plays a dominant role in infrared astronomy. A short review of atomic and molecular physics follows in order to provide a convenient reminder of the meaning of spectroscopic nomenclature. The second chapter covers the general properties of the Earth's atmosphere and outlines the main infrared surveys that have taken place or will shortly do so. Chapter 3 is devoted to photometry, emphasizing its fundamental aspects and the importance of traceability and calibration. Chapter 4 is an introduction to spectroscopy, treating firstly the stars. Photodissociation regions, which are of great interest in the infrared and millimeter region, are followed by HII regions and some “Rosetta Stone” spectra of representative objects. Chapter 5 is devoted to dust and its central role in star formation. Finally, because an understanding of the technology is important in obtaining reliable results, the last chapter covers the basics of infrared instrumentation.

Introduction

The central astronomical role of dust is at its most evident in the infrared. Protostars form from dusty clouds of molecular gas; the cool condensing dust around them emits copiously at sub-millimeter and far-infrared wavelengths. Even fully developed stars such as Vega may be found to be surrounded by remnant dust particles (see Habing et al., 1996), causing excess emission at long wavelengths. Near the end of their existence, a new generation of dust is formed by evolved stars, for example, in the atmospheres of asymptotic giant branch objects and in the ejecta of supernovae.

As in the visible region, dust scatters and absorbs light, giving rise to extinction, though its effects are much smaller in the infrared than in the visible. A dramatic example is the Galactic Center, which suffers 30 mag of extinction at V, so that only about one photon in 1012 comes through, whereas AK (2.2 μm) is about 2.5 mag and ∼10% of the photons can penetrate. Figure 5.1 shows the distribution of several types of objects with cool dust in an IRAS color-color diagram.

Dust also polarizes the light from distant stars and some properties of the polarization are found to be related to the extinction.

Even when the extinction is quite moderate, by observing in the infrared, the effects of interstellar absorption, which bedevil the use of the cepheid and RR Lyrae period-luminosity relations for distance determination, may be greatly reduced (e.g., Laney and Stobie, 1993, in the case of cepheids).

Introduction

We can consider spectroscopy at several different resolution levels.

The crudest resolution amounts to photometry, where the wavelength scale is divided into a small number of regions, or wide bands, to give an indication of a spectral energy distribution (SED). Examples might include the U(0.37), B(0.44), V(0.55), R(0.64) and I(0.80 μm) system in the visible region and the J(1.25), H(1.65), K (2.2), L(3.5), M(4.8), N(10), and Q(20 μm) system of ground-based infrared, as well as the 12, 25, 60 and 100 μm IRAS bands.

Many data with Δλ/λ ∼ 0.01 exist from spectrometers that used circular variable filters (CVFs). This is usually considered to be the least resolution which can really be called spectroscopy.

The next level might be called medium-resolution spectroscopy, where the detailed lines that usually compose molecular bands are not seen, but the bands are treated in an averaged manner. Most infrared spectroscopy falls into this category.

With high-resolution spectroscopy the individual molecular lines can be examined. Sensitivity considerations have limited this type of work hitherto to the study of the Sun, bright stars and planets. With the advent of large infrared array detectors and large telescopes, it is reasonable to predict that fainter objects will soon receive attention. High resolution has in the past been obtained through the use of Fabry–Perot etalons or Michelson interferometers, also called Fourier–transform spectrometers.

Stellar spectra

The theoretical study of stars is usually divided into stellar interiors and stellar atmospheres. The output of the interior can be regarded as something like a blackbody, but the atmosphere modifies the emergent flux according to its temperature structure, pressure structure and composition.

The birth of a new concept

The astronomical AO system developed at the University of Hawaii (UH) and its offspring, the AO user instrument of the Canada–France–Hawaii telescope (CFHT) are members of a new breed of AO systems based on the concept of wave-front curvature sensing and compensation (Roddier 1988). The concept emerged in the late 1980s at the Advanced Development Program (ADP) division of the US National Optical Astronomical Observatories (NOAO), as an output of a research program led by J. Beckers on the application of adaptive optics to astronomy.

Given the success of AO in defense applications, particularly surveillance systems, it was natural to seek components developed by the defense industry. The main difficulty was to obtain a good deformable mirror at a reasonable price. The technology being classified, one had no access to the latest developments. Commercially available mirrors were either monolithic mirrors, or first generation piezostack mirrors (see Chapter 4). Whereas monolithic mirrors of good optical quality could be purchased, their stroke was insufficient for the envisioned use on large astronomical telescopes. On the other hand, piezostack mirrors had enough stroke but were of poor optical quality and aged poorly. Most of all, the cost of a deformable mirror with suitable power supplies vastly exceeded budgets normally available for astronomical instrumentation.

It soon became clear that for surveillance applications one of the cost drivers was the high frequency response needed to follow the motion of satellites through the atmosphere.

An adaptive optics system (AOS) can be defined as a multi-variable servoloop system. Like classical servos, it is made of a sensor, the wave-front sensor (WFS), a control device, the real time computer (RTC), and a compensating device, the deformable mirror (DM). The goal of this servo is to compensate for an incoming optical wavefront distorted by atmospheric turbulence. It is designed to minimize the residual phase variance in the imaging path, i.e. to improve the overall telescope point-spread function (PSF). The input and the output of this servo are respectively the wave front phase perturbations and the residual phase after correction.

The design and the optimization of an AOS is a complex problem. It involves many scientific and engineering topics, such as understanding of atmospheric turbulence, image formation through turbulence, optics, mechanics, electronics, real time computers, and control theory. The goal of this chapter is to provide the basis of spatial and temporal controls. The reader will not find a tutorial on control theory but its application to an AOS. So the reader is assumed to be familiar with classical control theory (see for instance Franklin et al. 1990).

In the following, no restriction is made about the WFS or DM used. The principles are kept unchanged in the case of Shack–Hartmann, curvature sensor, or other WFS, and stacked array, bimorph, segmented DMs, or other wave front compensation devices.

In a first part, the control matrix determination and the modal control analysis of an AOS are described.

Introduction

Much of the early experimentation on astronomical adaptive optics was done in the 1970s on the Vacuum Tower Telescope (VTT) at Sacramento Peak (Buffington et al. 1977; Hardy 1981, 1987) either on stellar objects or on the sun itself. That telescope, although only 76 cm in aperture, was ideally suited for such experimentation because of its attractive environment for instrumentation. Diffraction limited imaging at visible wavelengths on both stars and the sun was achieved by Hardy. The solar results then clearly demonstrated the limitations on adaptive-optics-aided solar research resulting from the small isoplanatic patch size (a few arcseconds). Since then a few other efforts have been mounted to achieve diffraction limited imaging in solar observations.

Solar adaptive optics systems differ in a number of significant aspects from systems developed for night-time astronomy. Specifically:

1. (i) since the sun is an extended object, wave-front sensing on point-like objects as is done mostly in night-time adaptive optics systems is not an option.

2. (ii) solar seeing is generally worse than night-time seeing because the zenith angle/air mass at which the sun is being viewed is large in the early morning when night-time seeing still prevails and because seeing caused by ground heating by sunlight becomes severe later in the day when the sun is seen at greater elevations.

3. (iii) solar telescopes have generally much smaller apertures than night-time telescopes none, except for the 150-cm aperture McMath–Pierce facility on Kitt Peak, exceeding 1 meter in diameter.

4. […]