Preface

Progress in astronomy is fueled by new technical opportunities (Harwit, 1984). For a long time, steady and overall spectacular advances in the optical were made in telescopes and, more recently, in detectors. In the last 60 years, continued progress has been fueled by opening new spectral windows: radio, X-ray, infrared (IR), gamma ray. We haven't run out of possibilities: submillimeter, hard X-ray/gamma ray, cold IR telescopes, multi-conjugate adaptive optics, neutrinos, and gravitational waves are some of the remaining frontiers. To stay at the forefront requires that you be knowledgeable about new technical possibilities.

You will also need to maintain a broad perspective, an increasingly difficult challenge with the ongoing explosion of information. Much of the future progress in astronomy will come from combining insights in different spectral regions. Astronomy has become panchromatic. This is behind much of the push for Virtual Observatories and the development of data archives of many kinds. To make optimum use of all this information requires you to understand the capabilities and limitations of a broad range of instruments so you know the strengths and limitations of the data you are working with.

Introduction

Spectrometers divide the light centered at wavelength λ into narrow spectral ranges, Δλ. If the resolution R = λ/Δλ > 10, the goals of the observation are generally different from those in photometry, including both measuring spectral lines and characterizing broad features.

There are three basic ways of measuring light spectroscopically:

1. Differential-refraction-based, in which the variation of refractive index with wavelength of an optical material is used to separate the wavelengths, as in a prism spectrometer.

2. Interference-based, in which the light is divided so a phase-delay can be imposed on a portion of it. When the light is re-combined, interference among components is at different phases depending on the wavelength, allowing extraction of spectral information. The most widely used examples are diffraction grating, Fabry–Perot, and Fourier spectrometers. Heterodyne spectroscopy also falls into this category, but we will delay discussing it until we reach the submillimeter and radio regimes in Chapter 8.

3. Bolometrically, in which the signal is based on the energy of the absorbed photon. This method is applied in the X-ray, for example, using CCDs or bolometers, and will be discussed in Chapter 10.

Introduction

For many years, X-ray astronomy depended on gaseous detectors: basically, capacitors or series of capacitors with a voltage across them and filled with gas. Depending upon the value of the voltage, these devices: (1) just collect the charge freed when an energetic particle interacts with the gas (ionization chamber); or (2) provide gain [in order of increasing voltage, exciting the gas to produce ultraviolet light (scintillation proportional counter); creating modest-sized ionization avalanches to provide gain but with signals still in proportion to the original number of freed electrons (proportional counter); providing gain to saturation (Geiger counter); yielding a visible spark along the path of the avalanche (spark chamber)]. The absorbing gas is typically argon or xenon, for which high absorption efficiency in the 0.1–10 keV range requires a path of order 1 cm. Very thin windows are required to admit the X-rays to the sensitive volume – for example, 1 μm of polypropylene to provide > 80% transmission for energies above 0.9 keV. Proportional counters with multiple anode wires provide spatial resolution of a few hundreds of microns.

Because the atmosphere is opaque to them, X-rays and gamma rays require telescopes and detectors to operate from balloons, or more commonly from space (with the exception of the highest-energy gamma rays). Initially, the detectors were used without collecting optics; the large detector areas then resulted in high spurious detection rates due to cosmic rays. Anti-coincidence counters were required to identify charged particles coming from random directions, allowing probable X-ray events to be isolated. The necessity to operate at the top of or above the atmosphere plus these requirements on the detector systems were very limiting in terms of the angular resolution and sensitive areas that could be achieved.

Introduction

Most of what we know about astronomical sources comes from measuring their spectral energy distributions (SEDs) or from taking spectra. We can distinguish the two approaches in terms of the spectral resolution, defined as R = λ/Δλ, where λ is the wavelength of observation and Δλ is the range of wavelengths around λ that are combined into a single flux measurement. Photometry refers to the procedures for measuring or comparing SEDs and is typically obtained at R ~ 2–10. It is discussed in this chapter, while spectroscopy (with R ≥ 10) is described in the following one.

In the optical and near-infrared, nearly all the initial photometry was obtained on stars, whose SEDs are a related family of modified blackbodies with relative characteristics determined primarily by a small set of free parameters (e.g., temperature, reddening, composition, surface gravity). Useful comparisons among stars can be obtained relatively easily by defining a photometric system, which is a set of response bands for the [(telescope)-(instrument optics)-(defining optical filter)-(detector)] combination. Comparisons of measurements of stars with such a system, commonly called colors, can reveal their relative temperatures, reddening, and other parameters. Such comparisons are facilitated by defining a set of reference stars whose colors have been determined accurately and that can be used as transfer standards from one unknown star to another. This process is called classical stellar photometry. It does not require that the measurements be converted into physical units; all the results are relative to measurements of a network of stars. Instead, its validity depends on the stability of the photometric system and the accuracy with which it can be reproduced by other astronomers carrying out comparable measurements.

In 1934, two astronomers, Walter Baade and Fritz Zwicky, proposed the existence of a new form of star, the neutron star, which would be the end point of stellar evolution. They wrote:

These prophetic remarks seemed at the time to be beyond any possibility of actual observation, since a neutron star would be small, cold and inert, and would emit very little light. More than 30 years later the discovery of the pulsars, and the realisation a few months later that they were neutron stars, provided a totally unexpected verification of the proposal.

The physical conditions inside a neutron star are very different from laboratory experience. Densities up to 1014 g cm−3, and magnetic fields up to 1015 gauss (1011 tesla), are found in a star of solar mass but only about 20 kilometres in diameter. Again, predictions of these astonishing conditions were made before the discovery of pulsars. Oppenheimer & Volkoff in 1939 used a simple equation of state to predict the total mass, the density and the diameter; Hoyle, Narlikar & Wheeler in 1964 argued that a magnetic field of 1010 gauss might exist on a neutron star at the centre of the Crab Nebula; Pacini in 1967, just before the pulsar discovery, proposed that the rapid rotation of a highly magnetised neutron star might be the source of energy in the Crab Nebula.

On a time scale of some days, all pulsars show a remarkable uniformity of rotation rate. This is not surprising, since uniform rotation is exactly what is expected of a spinning body with a large stable moment of inertia and which is isolated in space. The angular momentum of the star can only change through the slowdown torque of the magnetic dipole radiation, or an associated material outflow, or, for the accretion-powered X-ray pulsars, the accelerating torque of in-falling material carrying angular momentum of a binary system. The effects on the radio pulsar are usually smooth and predictable: however, some very interesting irregularities in pulsar rotation have been observed, which are related to changes both within the interior of the neutron star and outside it.

These internal changes in pulsars appear to be spasmodic adjustments towards a slowly changing equilibrium state. For example, the rapidly rotating star will be appreciably oblate, and the equilibrium ellipticity will decrease during the slowdown; the crust is however extremely rigid and can only adjust to the changing ellipticity in a series of steps. The corresponding changes in moment of inertia might be large enough to be observable, since conservation of angular momentum will result in their being seen as steps in rotation rate. This effect is not in fact large enough to account for most of the observed repeated steps in rotation rate, and these are instead attributed to an interchange of angular momentum between the crust and the fluid interior.