The importance of the Sun as the most observable of all stars cannot be overstated. As shown in Figure 15.1, no other star can be studied with the degree of detail that we achieve in even the simplest observations of this source of all of our light and energy. As a result, what we have learned from the Sun we have applied in our study and analysis of the stars. Our knowledge of the sizes and distances of the stars is based upon our knowledge of the Sun. Also, we calibrate the luminosities of the stars in terms of our measurements of the output of energy from the Sun. In this chapter we shall first describe methods of observing the Sun in simple ways that can be used by anyone with a telescope. Then, we shall move on to more specialized methods and instruments that are used at observatories dedicated mainly to solar research.

Observing the Sun with a small telescope

The Sun is so bright that one should never try to make direct, naked-eye or telescopic observations of it. This is an absolute rule, for the observer can be blinded by even a brief attempt. There are, however, safe ways to view the Sun, and some of these require no complex equipment.

The most readily available method of seeing the Sun's apparent surface or photosphere is by means of eyepiece projection.

The general term “astrometry” is used to describe methods by which the positions of stars may be determined. We noted briefly in Chapter 1 that meridian telescopes were used to determine the right ascensions and declinations of stars, and we wish to repeat here that this work has been of fundamental importance to all astronomers. However, meridian telescopes have now been made obsolete by space-based observations and ground-based telescopes equipped with CCDs. Meridian telescopes were never an effective way to determine the coordinates of faint stars, galaxies, comets, and asteroids.

Today, most astrometric measurements are made with area detectors. The positions of stars on an image formed by a telescope are directly related to their actual positions in the sky, so it might seem that the analysis should be simple and straightforward. This is not really the case. The geometry in Figure 11.1 shows the basic relationships. The center of the lens of a telescope is at C and the focal plane is at FF′. A well-made lens should produce an image of plane GG′ in the plane FF′. The plane GG′ may be thought of as being tangent to the celestial sphere at point A. The sky appears as a spherical dome on which the stars appear as points. Thus, in a photograph a star at S is projected to T on the tangent plane, and an image of T is formed at T′.

Observational astronomers have taken advantage of each new development in detector technology. We compared visual observations, photographic plates, photoelectric photometers, and CCD cameras in Chapter 8. Our purpose in this chapter is to discuss how astronomers employ modern optical light detectors to measure astronomical sources in a scientifically useful way. We will focus our discussions on photoelectric and especially CCD photometry.

Fundamentally, there are two approaches to astronomical photometry. The simplest, differential photometry, compares sources sufficiently close together on the sky so that differential first-order extinction can be neglected. More complex is all-sky photometry, which takes full account of the extinction terms for stars observed far apart on the sky over the course of a night. In addition, photometry is done differently with photoelectric photometers compared to CCD cameras. We will begin with a short overview of photoelectric photometry and spend most of the chapter on CCD photometry.

Photoeletric photometry

As we noted in Chapter 8, a photoelectric photometer is of a fundamentally different design than a CCD camera. A single-channel photometer can be trained on only one star at a time. It is important that the star be carefully centered in its aperture. Often a range of aperture sizes is available; the smaller apertures are used in crowded fields and also when the seeing is very good. In between measurements of a star it is necessary to take measurements of the sky.

Many astronomy instructors adopted the first edition of Observational Astronomy as their primary text in advanced undergraduate courses. One of us (GG) used it as the primary text for an advanced undergraduate course in astronomy beginning in 1997. Unfortunately, the first edition went out of print in the late 1990s. By that time it had also become apparent that it was in need of revision. In particular, the charge coupled device (CCD) had already displaced nearly all other detectors in astronomy, but the first edition included only a short appendix on CCDs. Several chapters, instead, focused on photographic techniques. These included photometry, astrometry and spectroscopy. We have replaced all discussions of photographic techniques with CCD techniques in the present edition. We eliminated the chapters on classification of stellar spectra and radio astronomy and added chapters on light and detectors. In addition, we have reordered the material in several chapters in a way we hope is more pedagogically useful.

Most of the discussions about classical astronomical instruments, such as plate measuring engines and filar micrometers, have also been reduced or eliminated. The first edition remains a useful resource on these topics, and we encourage the interested reader to check with their local university library for copies.

The present edition of Observational Astronomy was the Master's thesis project of David Oesper at Iowa State University.

The intersection of three great circles on the surface of a sphere forms a three-sided figure. Such a figure is referred to as a spherical triangle, and it has some interesting properties. For example, the sum of the angles in a spherical triangle will usually be greater than 180°, whereas in a plane triangle this sum is exactly 180°. Consider the spherical triangle formed by the celestial equator and the hour circles of the vernal equinox and a star as in Figure 4.1. (An hour circle is a great circle that passes through a specified point on the celestial sphere as well as the north and south celestial poles.) The intersections of the three great circles have been labeled A, B, and C. Circles through the poles cross the equator at angles of 90°, so ∠B and ∠C are each equal to 90°, and the sum of the three angles will necessarily be greater than 180°.

The intersection of two planes cutting through a sphere forms an angle in a spherical triangle on the surface of the sphere. Each plane must pass through the center of the sphere. Thus, ∠A in Figure 4.1 is the angle between the hour circle planes through the vernal equinox and through the star. Note that ∠A is also equal to ∠BOC.

It is customary, just as in plane trigonometry, to label the sides of a triangle with lower case letters a, b, and c indicating the sides opposite the angles A, B, and C. Note that the length of a side is expressed in terms of an angle measured from the center of the sphere, as shown in Figure 4.2.

Today, astronomical databases once available only at the largest university libraries are but a mouse click away on the Internet. Except for a few very large catalogs described in Chapter 3, virtually any astronomical data can be downloaded from public-access websites. In addition, astronomical software has become more advanced and easier to use. Although it is not possible to give a complete listing of important astronomy software and Internet sites in this appendix, we have been careful to select those that should be most useful. While we cannot guarantee that all the websites listed below will remain active while this textbook remains in print, we have checked that they were active as of August 2005.

As stated earlier, the CCD has become the detector of choice in most astronomical applications. Its many advantages were listed at the end of the previous chapter, and some of its practical applications will be described in Chapters 10, 11 and 13. In this chapter we shall describe the necessary steps involved in eliminating noise and other sources of error so that CCD images can be used for accurate analysis.

Noise in the data

A raw CCD image contains information about your science targets. It also contains noise. Noise is usually categorized as random or systematic. Random noise causes a measured quantity to deviate from the “true” or “expected” value according to simple statistical relations, such as the normal distribution (see Appendix 1). Random noise cannot be eliminated; it can only be measured (characterized) so that its contribution to the signals from the science objects can be understood. A major advantage of a CCD over some other types of detectors is that its noise can be characterized accurately.

Systematic noise is caused by one or more processes that are not characterized by statistical distributions describing random events. Systematic noise can result from known sources or unknown ones. Known sources can be corrected; this is part of the calibration process. We will discuss below several types of known systematic noise in CCD data and how to correct for them. The presence of unknown systematic noise is always the greatest fear of the observational astronomer.

Students at the level of advanced undergraduates or beginning graduate students have often found that much information needed in the everyday practice of astronomy is not easily accessible. The necessary details are not to be expected in most textbooks, and one must often refer to early copies of some journals or to a professor's notes. It is my intention that this book should provide students with a ready reference of a practical nature.

For many years a course in astronomical techniques has been taught at Wellesley College, and the students there have been able to apply all of the methods described here. This book is thus based on the notes which I have developed while teaching this course. Over the years I have encountered a number of excellent books which were either to serve as texts for practical courses or as general handbooks for the use of amateur astronomers. My feeling has been that none of these covered the topics which I felt were most necessary at the level which I felt could be most useful.

It is my hope that this book will fill a real need in the reference material available to astronomers at many levels.

Time as we use it in our ordinary lives is based on the rotation and revolution of the Earth with respect to the Sun. These combined motions cause the Sun to appear to move continually around the heavens, and we define solar noon as the moment each day when the Sun crosses an observer's meridian. We could easily use the interval from one noon to the next to define the day, but in order to make all of the daylight hours part of the same calendar day, we start and end our days at midnight. The division of the day into twenty-four hours is strictly arbitrary and dates back to ancient Egypt, perhaps even earlier. The Greeks divided the periods of daylight and darkness into twelve equal parts to make twenty-four divisions in each day. The length of the hour defined in this way gradually changes throughout the year as the length of day and night varies with the seasons. The custom of dividing the hour into sixty minutes and the minute into sixty seconds is one of the last vestiges of the sexagesimal system of counting developed by the ancient Babylonians.

Today we keep track of time with a variety of clocks that range in complexity from sundials of the simplest sort to atomic clocks of the greatest precision.

We live on the surface of a planet that permits us to observe the distant Universe through a narrow window of the electromagnetic spectrum. While astronomers have expanded their vision with telescopes in Earth orbit, they still carry out most astronomical observations from the ground. As such, astronomers have to take into account the effects of the atmosphere on their observations. In this chapter we will discuss the five distinct ways the atmosphere affects the light from astronomical sources: extinction, refraction, seeing, scintillation, and dispersion.

Extinction

The Earth's atmosphere blocks most of the electromagnetic radiation arriving from space. Astronomers look through two “windows” in the electromagnetic spectrum to learn about extraterrestrial sources, the optical/near infrared (Figure 7.1) and radio. The optical window covers wavelengths from about 3000 Å to about 9000 Å. Several windows of moderately high transmission continue through the infrared up to about 26 microns. The windows in the radio band begin near 1 mm and end near 20 m. We will focus on the optical to near-IR window in this chapter.

We noted in Chapter 5 that the Johnson U bandpass is partly determined by the atmosphere's ultraviolet cutoff. Other bandpasses in the optical are determined by the filter transmission curves. The infrared JHKLM system (shown in Figure 7.1) is determined by the narrow windows between the deep molecular absorption bands.

Stars

When we wish to study the characteristics of a star, it is crucial that we be able to identify it with absolute certainty. Many resources are available to the astronomer to assist in star identification, and the method of choice depends largely on the brightness of the star. If the star is brighter than about fourth magnitude, the problem is quite simple, for there are not very many stars of comparable brightness. Once we set our telescope to the proper coordinates, there will probably be only one sufficiently bright star in the typical field of view. On the other hand, when we try to locate a star of eighth or ninth magnitude, there may be six or more stars in a field 30 arcmin in diameter. The numbers of progressively fainter stars in a given field increases dramatically. There are so many sources of error involved in setting a telescope that it is not reasonable to expect that the desired star will always be the one nearest the center of the field even if the coordinates are known and the telescope is set with great precision. Therefore, when trying to identify faint stars we must usually rely on maps or charts on which each star has been identified in some reliable manner. The observer can then compare the pattern of stars depicted on the chart with the pattern actually seen in the telescope and make a positive identification.

The development of spectacles (eyeglasses) was a tremendous impetus to the eventual development of the telescope. We do not know who first invented spectacles, but it is thought that they first appeared in Italy between 1285 and 1300. As eyeglasses became more common and more refined, it was perhaps inevitable that someone would figure out how to use lenses to form the first telescope. That happened in 1608, when the German-born Dutch eyeglass maker Hans Lippershey (1570–1619) built the first crude telescopes, not for astronomy, but for military uses.

In May of 1609, the Italian scientist Galileo Galilei (1564–1642) received word of this new magnifying instrument constructed using two lenses and a tube, and at once set about to build his own. In his Sidereus Nuncius (The Starry Messenger; 1610), Galileo describes his first telescope.

Galileo's first telescope could magnify objects just three times, and it was only about an inch in diameter. He soon constructed several more powerful telescopes, beginning at eight times magnification and eventually achieving 30 times.

What are astronomical spectra good for? For nearly 100 years photographs of low-dispersion spectra have been used to classify stars according to spectral type and luminosity class, determine spectral energy distributions, determine the redshifts and properties of galaxies, and measure emission line ratios in nebulae. High-dispersion spectra have been used to determine precise radial velocities, chemical element abundances in stellar atmospheres and in the interstellar medium, rotational velocities in stars, and magnetic field strengths in stellar atmospheres. We will explore aspects of some of these applications in this chapter.

Low-dispersion spectra are usually presented with absolute flux units on the vertical axis, as we showed in Chapter 5. High-dispersion stellar spectra, however, are usually normalized with respect to the local continuum. The continuum consists of the regions of highest flux that appear continuous except for the intervening absorption lines. Continuum windows are usually available even over a span of a few ångströms in high-dispersion spectra of Sun-like stars. In order to normalize a spectrum, first it is necessary to fit a low-order function through the continuum windows. This can be done by manually selecting the windows or by using a routine that automatically rejects high and low points according to an objective criterion. We show in Figure 13.1 a sample 50 Å region of a high-dispersion spectrum prior to and after continuum normalization. Normalized spectra are used in several of the analysis methods described below.

Visual observations

For thousands of years people recorded what they saw in the sky on rock walls, clay tablets, ivory and papyrus. In more recent times astronomers tried to reproduce on paper the precise patterns of the stars they observed with and without optical aid, some producing accurate and beautiful sky charts. Until the mid nineteenth century the human eye was the only available light detector.

The eye is truly a remarkable organ. Let us describe its structure and the function of some of its important parts in the overall process of vision. Figure 8.1 is a sketch of the right eye as it would appear looking down through the top of the head. The eye is essentially a spherical object that maintains its form by means of its tough outer layer, the sclera. The front center portion of the sclera is the transparent cornea through which all light entering the eye must pass. Behind the cornea is the crystalline lens, and the two are separated by a small amount of clear liquid known as the aqueous humor. The eyeball is filled with a jelly-like substance, the vitreous humor, which also helps it maintain its shape.

Light refracts at the outside surface of the cornea and at all of the interfaces within the eye. In a very real way, then, the optical characteristics of the eye are determined by the cornea, aqueous humor, crystalline lens and the vitreous humor.

One of the important actions that is in nearly all of the foregoing chapters is measurement. We measure time, coordinates, proper motions, parallaxes, magnitudes, the positions of lines in spectra, and shifts in positions of spectral lines, for example. After a measurement has been made we want to know how good the measurement really is, and in order to evaluate our measurements, we must turn to statistics. We would also like to use our measured data to make predictions either within or beyond the range of the measurements. Here we shall describe some of the principles that permit us to achieve these two goals in practical situations.

As an example, let us assume that a student is asked to determine the position of a spectrum line by measuring the distance from some reference position to the line center with a ruler. The smallest divisions on the ruler are one mm apart. We should be able to read the position of the line to within one tenth of a millimeter (0.1 mm). Just as a check the student makes a second setting and reads a new value with the ruler. She tries again and again until she has made fifty tries, and she never makes quite the same reading twice. Which of these many readings should be adopted as the correct one?