«W. Romanishin University of Oklahoma wjr The lastest version of this book (in pdf format) can be downloaded from ...»
An Introduction to Astronomical Photometry Using CCDs
University of Oklahoma
The lastest version of this book (in pdf format)
can be downloaded from http://observatory.ou.edu
This is version wrccd4a.pdf built on:
March 31, 2002
Foreword/ Thanks / Production Details
This book began as a set of lecture notes for a junior/senior course entitled “Observatory Methods” that I teach each spring at the University of Oklahoma (OU). The book is intended as an introduction for the college astrophysics major to photometry in the optical region of the spectrum of astronomical objects using CCD imaging from groundbased telescopes. Of course, in these times of Giga-buck satellite telescopes of various sorts, groundbased optical astronomy is only a part of observational astronomy. Within groundbased optical astronomy, spectroscopy, only brieﬂy mentioned here, probably takes up as much or more telescope time as photometry. That said, it is still obvious that imaging photometry is an important part of observational astronomy. With the ready availablity of inexpensive CCDs and computer power, even a small telescope can provide an important “hand on” learning experience not available with remote satellite observatories.
This book represents knowledge I have accumulated over 20 years of observing with a wide range of telescopes. My PhD disseration, ﬁnished in 1980, was probably one of the last observational dissertations to use photographic emulsions as the primary detector. Since then I have observed with various photomultiplier detectors and many diﬀerent CCD systems, on telescopes ranging in aperture from 0.4 to 10 meters.
I would like to thank the good people of the Great State of Oklahoma for paying my salary.
I would like to thank the NSF ILI program (Award Number 9452009) and OU for funds that purchased OUs 16 inch telescope and CCD. I would like to thank all the anonymous computer types who have written the free software used to produce this document.
This document was produced mostly using free software running under LINUX, with a little Windoze stuﬀ used only when unavoidable. ASCII LaTeX source text was edited with EMACS, an editor which looks exactly the same on my LINUX box in my oﬃce or on my ancient Windoze notebook while eating yet another bag of peanuts on the Southwest ﬂights to and from Arizona. My handrawn ﬁgures (and other ﬁgures swiped directly from other sources) were scanned with an HP 5200Cse scanner and saved as jpg images. The high resolution non- scanned plots were produced with IGI in STSDAS running under IRAF, saved as eps ﬁles. These eps ﬁles were converted to pdf ﬁles using epstopdf. The LaTeX ﬁles and jpg and pdf plots were converted into the ﬁnal pdf document with pdﬂatex.
Photometry: What and Why Many people are interested in astronomy because it is visually exciting. The many marvelous pictures of celestial objects taken using large telescopes on the ground or in space are certainly the most visible manifestation of modern research astronomy. However, to do real science, one needs far more than pictures. Pictures are needed as a ﬁrst step in classifying objects based on their appearance (morphology). To proceed past this initial stage of investigation, we need quantitative information- i.e. measurements of the properties of the objects. Observational astronomy becomes science only when we can start to answer questions quantitatively: how far away is that object?
How much energy does it emit? How hot is it?
The most fundamental information we can measure about celestial objects past our solar system is the amount of energy, in the form of electromagnetic radiation, that we receive from that object.
This quantity we will call the ﬂux. The science of measuring the ﬂux we receive from celestial objects is called photometry. As we will see, photometry usually refers to measurements of ﬂux over broad wavelength bands of radiation. Measurement of ﬂux, when coupled with some estimate of the distance to an object, can give us information on the total energy output of the object (luminosity), the objects temperature, and the objects size and other physical properties.
If we can measure the ﬂux in small wavelength intervals, we start to see that the ﬂux is often quite irregular on small wavelength scales. This is due to the interaction of light with the atoms and moleclues in the object. These bumps and wiggles in the ﬂux as a function of wavelength are like ﬁngerprints. They can tell us lots about the object- what it is made of, how the object is moving and rotating, the pressure and ionization of the material in the object etc. The observation of these bumps and wiggles is called spectroscopy. A combination of spectroscopy, meaning good wavelength resolution, and photometry, meaning good ﬂux calibration, is called spectrophotometry. Obviously, there is more information in a spectrophotometric scan of an object compared with photometry spanning the same wavelength range. Why would one do low wavelength resolution photometry rather than higher resolution spectrophotometry or spectroscopy, given the fact that a spectrum gives much more information than photometry? As we will see, it is much easier to make photometric observations of faint objects than it is to make spectroscopic observations of the same object. With any given telescope, one can always do photometry of much fainter objects than one can do spectroscopy of. On a practical note, the equipment required for CCD imaging photometry
is much simpler and cheaper than that needed for spectroscopy. With low cost CCDs now readily available, even small telescopes can do useful photometric observations, particularly monitoring variable objects.
Chapter 2 Visible EMR Almost all astronomical information from beyond the Solar System comes to us from some form of electromagnetic radiation (EMR). (Can you think of any sources of information from beyond the Solar system that do not involve EMR in some form?) We can now detect and study EMR over a range of wavelength or, equivalently, photon energy, covering a range of at least 1016 (ten thousand trillion sounds more impressive) - from short wavelength, high photon energy gamma rays to long wavelength low energy radio photons. Out of all this vast range of wavelengths, our eyes are sensitive to a tiny slice of wavelengths- roughly from 4500 to 6500 ˚. The range of wavelengths A our eyes are sensitive to is called the visible wavelength range. We will deﬁne a wavelength region reaching somewhat shorter (to about 3200 ˚) to somewhat longer (about 10,000 ˚) than the visible A A as the optical part of the spectrum. (Note: Physicists measure optical wavelengths in nanometers (nm). Astronomers tend to use ˚ngstroms. 1 ˚ = 10−10 m = 0.1 nm. Thus, a physicist would say A A the optical region is from 320 to 1000 nm.) All EMR comes is discete lumps called photons. A photon has a deﬁnite energy and frequency
or wavelength. The relation between photon energy (Eph ) and photon frequency (ν) is given by:
where h is Plancks constant and λ is the wavelength, and c is the speed of light. The energy of visible photons is around a few eV (electron volts). (An “electron volt” is a non- metric unit of energy that is a good size for measuring energies associated with changes of electron levels in atoms, and also for measuring energy of visble light photons. 1 eV = 1.602E−19 Joules.) The optical region of the spectrum, although only a tiny sliver of the complete EMR spectrum, is extremely important to astronomy for several reasons. Since our eyes are sensitive to this region, we have direct sensory experience with this region. Today, virtually no research level astronomical
12 CHAPTER 2. VISIBLE EMRobservations are made with the human eye as the primary detecting device. However, the fact that we see in visible light has driven a vast technological eﬀort over the past century or two to develop devices - photographic emulsions, photomultipliers, video cameras, solid state imagers- that detect and record visible light. The second overriding reason to study optical light is that the Earths atmosphere is at least partially transparent to this region of the spectrum- otherwise you couldn’t see the stars at night (or the Sun during the day)! Much of the EMR spectrum is blocked by the atmosphere, and can only be studied using telescopes placed above the atmosphere. Only in the optical and radio regions of the spctrum are there large atmospheric windows - portions of the EMR spectrum for which the atmosphere is at least partially transparent- which allow us to study the universe. Study of wavelengths that don’t penetrate the atmosphere using telescopes and detectors out in space- which we will call space astronomy - is an extremely important part of astronomy which has fantastically enriched our view of the universe over the past few decades.
However, space astronomy is very expensive and diﬃcult to carry out.
In purely astronomical terms, the optical portion of the spectrum is important because most stars and galaxies emit a signiﬁcant fraction of their energy in this part of the spectrum. (This is not true for objects signiﬁcantly colder than stars - e.g. planets, interstellar dust and molecular clouds, which emit in the infrared or at longer wavelengths - or signiﬁcantly hotter- e. g. ionized gas clouds, neutron stars, which emit in the ultraviolet and x-ray regions of the spectrum. Now, the next time you see the brillant planet Venus and think we are being invaded by space aliens, you may ask yourself why I included planets along with dust clouds in the above sentence. The reason is that the bright visible light you are seeing from Venus is reﬂected sunlight and not light emittedby Venus itself.) Another reason the optical region is important is that many molecules and atoms have electronic transitions in the optical wavelength region.
Imaging, Spectrophotometry andPhotometry
The goal of the observational astronomer to to make measurements of the EMR from celestial objects with as much detail, or ﬁnest resolution, possible. There are of course diﬀerent types of detail that we want to observe. These include angular detail, wavelength detail, and time detail.
The perfect astronomical observing system would tell us the amount of radiation, as a function of wavelength, from the entire sky in arbitrarily small angular slices. Such a system does not exist!
We are always limited in angular and wavelength coverage, and limited in resolution in angle and wavelength. If we want good information about the wavelength distribution of EMR from an object (spectroscopy or spectrophotometry) we have to give up angular detail. If we want good angular resolution over a wide area of sky (imaging) we usually have to give up wavelength resolution or coverage.
The ideal goal of spectrophotometry is to obtain the spectral energy distribution (SED) of celestial objects, or how the energy from the object is distributed in wavelength. We want to measure the amount of energy received by an observer outside the Earth’s atmosphere, per second, per unit area, per unit wavelength or frequency interval. Units of spectral ﬂux (in
cgs) look like:
(pronounced “f–nu”) if we measure per unit frequency interval.
Figure 3.1 shows a typical spectrum of an astronomical object.
This covers, of course, only a very limited part of the total EMR spectrum. Note the units on the axes. From the wavelength
14 CHAPTER 3. IMAGING, SPECTROPHOTOMETRY AND PHOTOMETRYcovered, which lies in the UV (ultraviolet), a region of the spectrum to which the atmosphere is opaque, you can tell the spectrum was not taken with a groundbased telescope.
fλ and fν of the same source at the same wavelength are vastly diﬀerent numbers. This is because a change of 1 ˚ in wavelength corresponds to a much bigger fractional spectral coverage A
than a change of one Hz in frequency, at least in the optical. The relationship between fλ and fν is:
c fλ = fν (3.3) λ2 Spectrophotometry can be characterized by the wavelength (or frequency) resolution- this is just the smallest bin for which we have information. E.G. if we have “1 ˚” resolution then we know A ˚ngstrom interval.
the ﬂux at each and every A We characterize the wavelength resolution by a number called the “resolution”:- this is the wavelength (λ) divided by the wavelength resolution(∆λ). E.G. If the wavelength resolution element is 2 ˚, and the observing wavelength is 5000 ˚, then the resolution is 2500.
A A To get true spectrophotometry, we must use some sort of dispersing element (diﬀraction grating or prism) that spreads the light out in wavelength, so that we can measure the amount of light in small wavelength intervals. Now, this obviously dilutes the light. Thus, compared to imaging, spectrophotometry requires a larger telescope or is limited to relatively bright objects.
Spectrophotometry also requires a spectrograph, a piece of equipment to spread out the light.
Good research grade spectrographs are complicated and expensive pieces of equipment.