ASTR 5110, Majewski [FALL 2017]. Lecture Notes

## DETECTION LIMITS AND THE ARRAY EQUATION

 REFERENCES: Howell, Handbook of CCD Astronomy, Section 4.4. Newberry 1991, PASP, 103, 122. Gullixon 1992, ASPC, 23, 130.

The Array Equation: Simple Photometric Case

In planning and carrying out an observational experiment, it is important to have an understanding of the needed, expected, and actual final S/N of the experiment.

• The S/N is a measurement of the quality of the experimental result.

• You should have a thorough understanding of how to calculate this S/N for any instrument you use.

• Section 4.4 of Howell is a must read for a beginning understanding on noise in CCD arrays.

• The discussion below follows Howell's description closely.

Previously we discussed sources of noise encountered when working with a detector array, and considered two extreme limits:

• the readout-limited regime, when the noise is dominated by the readout noise of the detector,
• the sky-limited regime, when noise from the sky background dominates.

Now we want to be more explicit and carefully calculate all contributions to the signal and noise when using a detector array like a CCD.

• Though we are discussing the CCD explicitly here, the equations below generally work for many other kinds of array detectors as well.

In the case of imaging with an electronic array detector, like a CCD, and a source extending over npix pixels, the S/N of the detection of this source can be estimated in the most basic way as follows:

• It is important to note that the statistics of counting depends on the number of photons converted to electrons by the detector and NOT on their ADU representation in the CCD image.

• Thus, all quantities in the above equation are in units of electrons, which can be obtained from the quantities in ADUs by knowing the inverse gain (electrons/ADU).

• The other terms in the denominator can be understood in terms of Poisson statistics (see below).

Clearly, the signal is given by N*, the total photoelectrons obtained from the source and collected in our "measurement".

• This "measurement" may be from one pixel (if measuring the S/N per pixel -- as is sometimes done in spectroscopy or in measures of a surface brightness).

• Or the "measurement" may be from multiple pixels, such as all those contained within a stellar profile or within an aperture.

The random fluctuations in the count, or "noise" N, from the various contributed sources are uncorrelated and therefore combine in quadrature.

The total noise is then given by:

where the terms in the square root are, in succession:

• The Poisson noise in the source counts, (N*)1/2.

• The Poisson noise in the sky background counts, (NS)1/2, with equal contributions (added in quadrature) from npix pixels in the "measurement".

• The Poisson noise in the dark current counts, (ND)1/2, with equal contributions (added in quadrature) from npix pixels in the "measurement".

• The contribution of read noise, which is not a Poisson noise and is usually characterized as NR electrons per pixel (so, in the quadrature addition, appears as a squared quantity).

As may be seen, for bright sources, the first source of noise dominates and we have the classical

For fainter sources, the other noise terms become important (e.g., a read-noise-limited measurement has the errors dominated by NR).

The Array Equation: More Exact Case

In some cases (high backgrounds, faint sources, poor spatial sampling, large inverse gains -- see Howell) a more complex form of the equation is required:

where

• ( 1 + npix / nB ) is a coefficient that accounts for noise incurred due to an inaccurate estimation of the background level on the CCD image (where nB is the number of pixels used to measure it).

Clearly nB being large is better.

• The term involving the inverse gain, G in electrons/ADU, accounts for digitization noise within the A/D converter, an error that becomes considerable as the inverse gain grows.

σ2f is the 1σ error introduced in the A/D converter of the CCD when it has to decide between two ADU levels bracketing the electron count divided by G.

The value of this quantity is taken to be σ2f = 0.289.

Standard Error in Magnitudes

As we have said before, and will explore in more detail soon, the magnitude system is defined on a logarithmic scale, where m = -2.5 log10 F + constant, where F is the flux.

If one recognizes that the standard error (standard deviation) of a measurement, σ, is related to the S/N as

then we can derive the error in magnitudes as

where p is

and the factor 1.0857 is the conversion factor one obtains when one calculates the error in magnitudes (not electrons -- derive this yourself!).

Required Integration Times

One of the most important things one needs to calculate in preparation for an observing run is how long one should expect to integrate to achieve a given S/N.

• Many observatories have exposure time calculators that will derive these values for you for a given instrument.

• At minimum, these observatories should be able to tell you either directly or give you the information needed to find the count rate expected for a source of a given magnitude for that instrument.

Then, if instead of thinking of counts, we have the above equations expressed with count rates per second, we can rewrite the S/N in terms of integration time, t, as

One then solves this quadratic equation for t to know how long to integrate to obtain a given S/N within a detection area npix (see Howell for an explicit solution).

The main thing to notice is that in the source-limited regime,

S/N ∝ t1/2.

The Array Equation: Spectroscopic Case

In spectroscopic applications, it is common to quote the S/N per pixel or S/N per resolution element in the continuum.

• In these cases, one must adjust npix accordingly, taking into account both the number of pixels in the dispersion direction (e.g., per resolution element) and those summed perpendicular to the dispersion direction in the extraction of the source spectrum.

One can also quote the S/N for a given spectral line or feature.

Note that throughputs of spectrographs are much reduced compared to direct cameras, amounting generally to a few to 10% or so, even discounting the possibility of light lost by not making it through the slit.

• Consider the additional amount of optics between the telescope focal plane and the CCD chip in the case of spectroscopy.

AAO photographs from http://www.aao.gov.au/2df/gallery.html. All other material copyright © 2002, 2005, 2011, 2015, 2017 Steven R. Majewski. All rights reserved. These notes are intended for the private, noncommercial use of students enrolled in Astronomy 511 at the University of Virginia.