
\chapter{CCD Detectors} 
\section{Introduction}

The two types of electronic detector available with the IDS are the Image
Photon Counting System (IPCS), and a range of Charge Coupled Devices (CCDs). If
you are reading this manual you have already decided that the CCD is best
suited for your particular observation. The Observer's Guide has a thorough
discussion comparing both types of electronic detector available with the IDS
(CCD and IPCS); see also Gemini \# 21. We will include here the relevant
conclusions for a CCD user. The CCD is described in Sections 2.2 and 2.3. (If
you are interested in the IPCS, see the ``IPCS  Cook Book'' and the ``IDS +
IPCS'' Manual, currently in preparation). 

\section{Overview}

 A CCD is a 2-dimensional electronic  detector. It has enormous attractions for
astronomers: linearity, with dynamic ranges of 10000 being realized;
photometric accuracy, better than 1 per cent; speed,  with  quantum
efficiencies of up to 80 per cent; wide wavelength response, some  devices
being usable from the near UV ($<$ 300 nm) to the near IR  ($>$ 1 $\mu$m) and
digital output, with data readily available in the computer. In addition, it is
mechanically stable, extremely tough and difficult to damage by excess
illumination.  The main disadvantages of CCDs (as available at present at the
INT) are the small size as compared to photographic  plates or an IPCS, and the
presence of readout noise, which renders it  less suitable than an IPCS for the
study of very faint or rapidly varying objects. 

This section is a brief  description of a CCD. Some useful references for
further  reading are: 

\hspace*{1.0cm} {\it ESO-OHP workshop on the Optimization of the Use of CCD
detectors \\* 
\hspace*{1.0cm} in astronomy, eds. Baluteau, J.P. \& D'Odorico, S., 1986} \\*
\\*
\hspace*{1.0cm} Wall, J.V. \& Laing, R.A. {\it User's Guide to the prime focus
CCD camera on the INT} \\* 
\\*
\hspace*{1.0cm} Argyle, R.W., Mayer, C.J., Pike, C.D. \& Jorden, P.R. \\*
\hspace*{1.0cm} {\it User Guide to the JKT CCD camera, 1988.} \\*
\\*
\hspace*{1.0cm} Jorden, P.R. {\it Basic parameters of CCDs in use at La Palma,
\\* 
\hspace*{1.0cm} La Palma Technical Note 55, revised September 1988} \\

\section{General description}

 A CCD is a silicon based semiconductor arranged as an array of photosensitive 
elements, each one of which generates photoelectrons and stores them (to use
the now familiar figure of speech) as a  small `bucket' of charge.  Each pixel
is typically  20 to 30 $\mu$m square.  CCDs in use as astronomical detectors
for the  past five years have about 200,000 pixels, usually arranged in a
format of  about 500 rows by 400 columns, and the active area is thus about 1
square  centimetre. 

 When requested, the elements form a bucket brigade;  each row of charges is
passed from element to element, a process which is  known as clocking, down the
columns and horizontally along the final row.  The value in each pixel is
measured in turn and recorded digitally.  To ensure that only positive numbers
result from this analogue to  digital conversion process a fixed offset known
as the bias level  is introduced. The charge-transfer process is essentially
noise-free  and almost all of the noise contributed to the signal by the CCD is
 from the output stage, where the charge content of each bucket is  measured. 
This is called the readout noise. 

 The dark current of CCDs means that they must be cooled to cryogenic 
temperatures for use as astronomical detectors.  At room temperatures,  the
dark signal is such that most CCDs would saturate  in $<$ 1 second.  Cooling
via temperature-controlled liquid-nitrogen dewars  or by closed-cycle
refrigeration provides a crucial reduction in dark  current.  The mobility of
electrons is somewhat impaired by cooling, so  that a compromise is required to
maintain adequate charge-transfer  efficiency.  The optimum temperature is
about 150 K. 

\section{Design} 

 CCDs require electronics to generate electrode biasses, clocking waveforms,
low-noise amplification of weak signals, A/D conversion, buffering, and digital
 storage.  They also require temperature-controlled environments mountable at 
different telescope foci for direct imaging, and at spectrograph foci for use 
as spectroscopic detectors.  The details of how operational systems meet these 
demands differ, but the major elements are similar. 
 
 The cryostat or dewar provides the cooled environment for the CCD. A
liquid-nitrogen cryostat designed at RGO can be used in either downward or
upward-looking modes, for example in direct imaging at prime or Cassegrain
foci.  It  has a hold time of some 12 hours so that it only has to be filled 
once a night.  The CCD itself is mounted behind an anti-reflection  coated
quartz window, on a copper block with which it makes good thermal  contact and
which has a temperature sensor attached.  Resistive heating controlled by a
feedback loop maintains the CCD  at the optimum operating temperature of 150 K
to within $\pm$0.05 K.  A preamplifier is mounted on the cryostat. 

The main functions  of the electronics needed to drive the CCDs are: analogue
signal sampling  and CCD control, digitisation of data,  telemetry of system
voltages, control interfaces and data links  to the instrument computer. The
systems may be reprogrammed in order  to use a wide variety of different chip
types or operating modes.  Windowed readout of the chip can be used to minimise
readout time  or data storage. On-chip binning (available for most chips) can
be used to  increase signal to noise at the expense of spatial resolution. 

\section{CCD systems on the INT} 

 The CCD systems on the INT use either an RCA type SID501 chip,  or a GEC
P8603,  dye coated for improved blue response. Table 2.1 summarises  important
general parameters of the two types of CCD chip available, and  has been taken
with some modification  from La Palma Technical Note 55, which describes the
CCD systems available  in rather more detail. 
 
 It can be seen that the two chips complement each other well. The RCA SID501
is larger and has a higher efficiency at most wavelengths, and so is better
suited for those imaging applications requiring a large field and deep,
broad-band exposures. The GEC P8603 has a much lower readout noise and  so will
be better suited for most spectroscopic applications. The principal limitation
of the RCA chip for spectroscopy is that it is thinned and hence gives rise to
very large amplitude interference fringes. {\bf Hence, only the GEC chip is
offered with the IDS.} 


\begin{table}[htp]
\caption{CCD parameters}
\begin{center}
\begin{tabular}{lcc}
\\*
 Chip                   &         GEC P8603    &     RCA SID501      \\
 Pixel size ($\mu$m)    &   22 $\times$ 22     &    30 $\times$ 30    \\
 Format (pixels)        &  385 $\times$ 578    &   320 $\times$ 512    \\
 Format (mm)            &  8.5 $\times$ 12.7   &   9.6 $\times$ 15.4   \\     
 Readout frame (pixels) &  400 $\times$ 590    &   350 $\times$ 512   \\
 Readout time (seconds) &      $\sim$25        &       $\sim$10   \\
 Electrons per ADU      &       $\sim$1        &        $\sim$4   \\
 Saturation (ADU)       &      $>$60000        &      $>$60000   \\
 R.Q.E. (per cent)      \\
\hspace*{0.5cm} Peak    &             55       &            75   \\
\hspace*{0.5cm} 400 nm  &             15       &            40   \\
\hspace*{0.5cm} 800 nm  &             40       &            60   \\
Noise:                  \\
\hspace*{0.5cm} Readout noise      &   5--10      &      $\sim$60 \\
\hspace*{0.5cm}  Cosmic ray events &   1--2       &      $\sim$6   \\
\hspace*{0.75cm} (events frame$^{-1}$ min$^{-1}$) \\
\hspace*{0.5cm}  Dark count        &   $<$1       &       $<$1   \\
\hspace*{0.75cm} (ADUs pixel$^{-1}$ hour$^{-1}$)  
\end{tabular}
\end{center}
ADU=Analogue-Digital conversion Units.\hfill\break
R.Q.E.=Responsive Quantum Efficiency.
\end{table}

\subsection{Performance}


\subsubsection{Efficiency} 

A major advantage of CCDs is their high efficiency. The nominal wavelength
response for the two types of CCD currently in use on La Palma is shown in
Figure 2.1. The RCA SID501 has a peak    efficiency of about 75 per cent,
whilst the GEC P8603 has a peak efficiency    of about 55 per cent. Both chips
have a useful sensitivity over most    of the optical/near IR spectrum. Note
however that although the useful    response of the GEC P8603 extends down to
the atmospheric cutoff    at about 3000 \AA, the response of the RCA SID501
cuts off sharply below about 3500 \AA.  Our measurements however indicate that
the blue performance of the RCA is significantly worse than this, and is
completely useless in the U band. The GEC response is close to that shown. 


\begin{figure}[p]
\vskip 20cm
\caption[CCD efficiency curves]
{Responsive quantum efficiency of the CCD chips available in the INT}
\end{figure}

\subsubsection{Noise level} 

   The read out noise for the CCD chips used on La Palma    is given in Table
2.1. This is typically 5-10 electrons per pixel for the    GEC P8603 and about
60 electrons per pixel for the RCA SID501 

   Cosmic ray events become important for longer exposures.    During an
exposure cosmic rays hit the CCD.  When a pixel is    hit, its charge content
is vastly increased and the resulting `spike'    in the image needs to be
removed during the data reduction stage.     This can be done either by
numerical    filtering techniques using the fact that the appearance of a
cosmic ray    event is quite distinct from that of an astronomical object, or
by    combining multiple exposures of the same object.  The former process is
likely to be interactive and (therefore slow) and a ``patch'' with information
in the strike zone is irretrievably lost. The latter is recommended. 

\subsubsection{Saturation level and dynamic range} 

   The peak capacity of a CCD pixel    is typically 200,000 electrons, and CCDs
are demonstrably linear devices    up to within a factor of two of this limit. 
When the limit is    exceeded, the arrival of more photons continues to create
further    photoelectrons, which then spread out along the column of the CCD. 

   In addition to saturating the chip, it is possible to saturate    the
analogue to digital converter. Although the maximum digitised    count is
65,535 (2$^{16}$-1), it is advisable to keep the number    of counts below
about 35,000 to avoid non-linear effects. 

   Together with readout noise, the saturation level defines the dynamic range
of a CCD, crudely, as (saturation level)/(readout noise). 

\subsubsection{Charge transfer efficiency} 

At low light levels the charge transfer is sometimes imperfect and the bucket
brigade spills some electrons along the way during read-out.  This results in
slightly deformed images, showing a trailing along the rows and/or columns of
the CCD.  Faint asymmetric wings may be generated around bright emission lines.
To improve the performance one can preflash the chip by pre-illumination with a
fixed amount of light.  This is, of course, undesirable because it adds shot
noise to the data.  Although all the chips currently in use on La Palma have
good charge-transfer characteristics and are said by the manufacturers not to
require preflash, the charge-transfer is not perfect. Therefore we recommend
that users consult their support astronomer over whether a preflash is needed. 

\subsubsection{Cosmetic defects} 

    Cosmetic defects can come in many varieties, and vary    from chip to chip.
Dark columns, areas where charge-transfer is inhibited    unless a threshold of
electrons is present, can be dealt with effectively    by pre-flashing. 
Defective areas such as bad columns, bad pixels or hot    pixels cannot be
cured. Each chip has its ``fingerprint'' of cosmetic    defects and care should
be taken to avoid the worst affected areas when    positioning an object on the
chip.  Multiple exposures in which the    object of interest is shifted to a
different location on the chip between    exposures can be of great help to
obtain clean images. 

Other factors in addition to the required signal to noise ratio which need to
be taken into account when choosing a detector are: 


\subsubsection{Detector size} 

For certain types of experiment, it may be important to note that    the IPCS
is a much larger detector than any of the CCDs currently    available on the
INT. For spectroscopic observations, the INT IPCS    provides 2044 pixels along
the dispersion direction, each    about 15 $\mu$m in size. The GEC CCD   
provides 578 pixels along the dispersion direction, each 22$\mu$m    in size.
This means that for a fixed wavelength resolution,    the IPCS provides almost
2.5 times the wavelength coverage.    The usefulness of this multiplex
advantage depends very much    on the type of experiment being performed. For
an experiment    where wavelength coverage is important, it might mean that   
the IPCS is the preferred detector, even though the integration    time to
reach the desired signal to noise ratio in    each resolution element is 2-3
times longer. A new generation of CCDs    with $1000^2$ to $2000^2$ pixels is
becoming available for astronomical observations and we hope to offer it on the
INT from 1990. 
   
\subsubsection{Real time data display} 

   The data cannot be inspected in real time with the CCD. This is a serious
drawback when doing surveys, because one cannot terminate an exposure on the
basis of the actual    signal to noise ratio reached, as is the case with the
IPCS. 

\subsubsection{Saturation} 

The exposure time for bright objects could be limited by saturation. 

\subsubsection{Cosmic ray events} 

   These are an important source of additional noise on a CCD, limiting in
practice the exposure time. Cosmic rays which strike the detector produce
bright and narrow noise spikes in the signal. The effect on the data depends
exactly on where on the detector the event occurs. The best way to deal with   
cosmic rays is to break long exposures into several shorter ones. This has the
major disadvantage of increasing the total    readout noise. If the total
number of events is small they can usually be succesfully removed by software. 

\section{Exposure times}


Once the signal to noise ratio required for the observation to be successful
has been defined, you want to estimate the integration time needed to achieve
it. Here we outline how to perform this calculation. 

The noise level for a typical astronomical observation will consist of three
components. The first component is signal independent noise arising from the
detector itself (readout noise). Secondly, we have shot noise in the signal
from the object being observed. This follows a Poisson distribution, so its
variance is equal to the signal level. Finally, we also have shot noise in any
background signal (particularly the sky background and preflash). Adding the
noise terms in quadrature, the signal to noise ratio ($SNR$) is given by: 

\begin{equation}
 SNR = \frac{ E \times N_{obj} \times t } 
   {\sqrt{ ( E \times N_{obj} \times t ) + ( E \times N_{back} \times t ) 
   + \sigma^2 } }
\end{equation}

\begin{tabbing}
aaaaa\=aaaaaa\=aaa\=  \kill
\>$E$ \>= \>Responsive quantum efficiency of the detector\\*
\>$N_{obj}$  \> = \> Number of photons per resolution element per second incident on detector \\*
\>          \>    \>from object \\*
\>$N_{back}$  \> = \> Number of photons per resolution element per second incident on detector \\*
\>          \>      \> from background (e.g. sky) \\*
\>$t$ \>= \> Integration time \\*
\>$\sigma$ \>= \>Readout noise \\
\end{tabbing}

The signal to noise ratio obtained therefore depends on both the efficiency of
the detector and the level of detector noise. 

Table 2.2 gives estimates of the efficiency of the GEC CCD at a number of
wavelengths from Figure~2.1. 

\begin{table}[htp]
\caption{CCD efficiency}
\begin{center}
\begin{tabular}{cc}
Wavelength &  Detector efficiency (\%) \\
   (\AA)     &  GEC P8603  \\
\\                                             
  3500    &    12  \\
  4000    &    15  \\
  4500    &    15  \\
  5000    &    24  \\
  5500    &    35  \\
  6000    &    42  \\
  6500    &    50  \\
  7000    &    50  \\
  7500    &    51  \\
  8000    &    45  \\
  8500    &    33  \\
  9000    &    22  \\
  9500    &    10  \\
\end{tabular}
\end{center}
\end{table}

The detector noise of a CCD is dominated by the readout noise, which is 5--10 
electrons per pixel for the coated GEC P8603 (Table 2.1). Note however that the
number we really need is the readout noise per resolution element, not the
readout noise per pixel. The readout noise per resolution element depends on
two factors: 

\begin{itemize}
\item How many CCD pixels are there in a resolution element ? In 
    slit spectroscopy of a point object, the emission 
    might easily be extended for several pixels along the slit, and we 
    will need to bin these up in order to obtain a spectrum. Similarly, 
    for an experiment to detect a well-resolved emission line, it may 
    be possible to bin up several
    pixels in the dispersion direction into one resolution element.
    Note that according to sampling theory, a resolution element
    must always be at least two pixels in each direction.         
    In order to obtain the effective readout noise per resolution
    element, the readout noise per pixel must be multiplied by the
    square root of the number of pixels in a resolution element.

\item How many separate CCD exposures are being combined?
    The target object may be so faint that a very long integration time 
    is necessary to obtain the desired signal level. Figures 2.3 to 2.5 
    can be used to estimate the total integration time required. 
    If the total integration time is greater than about 1/2 an hour, it 
    will be necessary to split the observation into a number of shorter 
    exposures. Each CCD exposure adds in more 
    readout noise, resulting in the effective detector noise being equal 
    to the readout noise per exposure times the square root of the number 
    of exposures. 
\end{itemize}

A first relatively straightforward but important case is where the noise level
is dominated by shot noise on the sky background (sky limited observations).
This is usually the case for low resolution spectroscopy or imaging
observations. Typical values for the sky brightness on La  Palma in a variety
of wavebands are given in the Observer's Guide. As far as signal to noise is
concerned, the most effective detector for sky limited observations is simply
the most efficient. 

We now consider the case where the sky background can be neglected, normally
the case for high resolution spectroscopy or bright objects. As we noted above,
observations at high signal levels, where we might expect a high signal to
noise ratio, will be best suited for the CCD, whereas observations at low
signal levels, and low signal to noise ratio, will be best suited for the IPCS.
The basic message here is that if you intend to take short exposures of fairly
bright objects, you should probably use the CCD, whereas if you expect to
observe the same object all night to obtain adequate signal to noise ratio (S/N
$<$ 30), you should probably be using the IPCS. 

As discussed above, the signal to noise ratio obtained from an observation
depends on two factors: the level of detector noise, and the number of photons
detected in each resolution element. In what follows we present the data
necessary to estimate the number of photons incident on the detector in each
resolution element. 

For a system with no light losses, it is straightforward to show that the
number of photons incident in each resolution element is given by: 
  
\begin{equation}  
N = 15092 \times A \times F \times f \times \frac{\delta \lambda}{\lambda} 
\end{equation}

\begin{tabbing}
aaa\=aa\= \kill
$N$ \> = \> Photon flux at the detector per resolution element (photons sec$^{-1}$) \\
$A$ \> = \> Geometric aperture of the INT telescope (m$^2$) \\
     \> = \> 4.41 \\
$F$ \> = \> Monochromatic flux density of object (mJy) \\
    \>   \>  If you like to work in magnitudes you should use the fact that \\
    \>   \>  an object with a magnitude of m$_{AB}$ = 15 has a flux density of 3.6 mJy. \\
$f$ \> = \> Fraction of total flux density of object in each spatial resolution element. \\
   \>    \> Note that for slit spectroscopy, the spatial resolution perpendicular to the \\
  \>    \>  direction of the slit is set by the slitwidth. For many observations, the  \\
  \>    \>  slitwidth will be comparable to the size of the seeing disc, and so even for \\
  \>    \>  sources which are intrinsically unresolved, a significant fraction of the light \\ 
  \>    \>  will not enter the spectrograph. \\
$\delta \lambda$ \> = \> Observing bandpass (\AA). For imaging observations, this is the FWHM of the filter  \\
  \>    \>  used. For spectroscopy, it is equal to the wavelength resolution of the observations. \\
$\lambda$ \> = \> Observing wavelength (\AA)
\end{tabbing}

Unfortunately, only a small fraction of these photons are actually detected,
perhaps 1 per cent! This low efficiency is due to a number of factors,
discussed in what follows. 

\subsection{Atmospheric extinction} 

   La Palma Technical Note 31 (King, 1984) discusses the atmospheric   
extinction on La Palma in some detail. Briefly, it is possible to separate   
the extinction into a wavelength dependent component, due to Rayleigh   
scattering by air molecules and absorption by ozone, and a component due to
dust scattering which is to some extent wavelength independent because large
dust particles dominate the size distribution. The wavelength dependent   
component is plotted as a function of wavelength and airmass in Figure 2.2.
Table 2.3 gives the airmass as a function of zenith distance.    It can be seen
that this component of the atmospheric extinction is    much more important in
the blue than    in the red, varying from 3.7 magnitudes per unit airmass at
3000 \AA\ to 0.007    magnitudes per unit airmass at 10000 \AA. 

\begin{figure}[p]
\vskip 20cm
\caption[Atmospheric extinction on La Palma]
{Theoretical atmospheric extinction on La Palma as a function
of wavelength and airmass}
\end{figure}

   The extinction due to dust scattering    varies from night to night, but is
usually less than a few tenths of a magnitude per unit airmass. The total
vertical extinction in the V band (i.e.    the sum of the wavelength
independent and wavelength dependent components)    is measured each night by
the Carlsberg Automatic Meridian Circle, and users    with access to STARLINK
may examine these results in the file    {\ss
RGVAD::SYS\$SYSDEVICE:[LPINFO.EXTINCTION]CAMCEXT.DAT.}. These    data may be
used to estimate the dust extinction on any particular night. Obviously, the
use of a plane-parallel azimuth-independent geometry (the sec~z law) is less
accurate for dust extinction than for air. 

\begin{table}
\begin{center}
\caption{Airmass as a function of zenith distance}
\begin{tabular}{cc}
Zenith distance  & Airmass \\
(degrees)    [Ou    &         \\
0                & 1.000   \\
20               & 1.064   \\   
40               & 1.304   \\
50               & 1.553   \\
55               &  1.740    \\
60               &  1.995    \\
65               &  2.357    \\
\end{tabular}
\end{center}
\end{table}


\subsection{Telescope efficiency} 

   Light is lost at each reflection from an aluminised    telescope mirror. It
is difficult to quantify this precisely, since    the reflectivity of the
mirror degrades with time, but a value    of 85 per cent for the reflectivity
would be a reasonable estimate. 

\subsection{Spectrograph efficiency} 

   The efficiency of the INT IDS is best considered as the product of three   
components; collimator, grating and camera. The efficiencies of the    3
collimators available are plotted in Figure 2.3, and the efficiencies    of the
various gratings are plotted in Figure 2.4. The efficiencies    of the 2
cameras are discussed in Chapter~3. 

\begin{figure}[p]
\vskip 20cm
\caption[Collimator efficiencies]{Efficiency curves for the IDS collimators}
\end{figure}

\begin{figure}[p]
\vskip 20cm
\caption[Grating efficiencies]{Grating efficiencies as a function of wavelength}
\end{figure}
\begin{figure}[p]
\addtocounter{figure}{-1}
\vskip 20cm
\caption[(Cont.)]{(Cont.)}
\end{figure}
\begin{figure}[p]
\addtocounter{figure}{-1}
\vskip 20cm
\caption[(Cont.)]{(Cont.)}
\end{figure}

\subsection{Filters} 

   Filters will be used in spectroscopic    observations, when for example it
is necessary to reject light    from unwanted orders or spectral regions. 

   Figure 2.5 shows the wavelength    dependence of transmission of the Schott
glass filters used on    La Palma. 


\begin{figure}[p]
\vskip 20cm
\caption[Transmittance of Schott glass filters]{Transmittance of Schott 
glass filters as a function of wavelength}
\end{figure}

\subsection{Detector efficiency} 

The efficiency of the two types of CCD used on La Palma is shown in Figure 2.1.








%\begin{figure}[p]
%\vskip 2cm
%\caption[Transmittance of the copper sulphate filter]
%{Transmittance of the copper sulphate filter as a function
%of wavelength}
%\end{figure}

\vfill

