Introduction to CCDs

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Introduction to CCDs

Basic principles of CCD imaging

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What is a CCD ?

Charge Coupled Devices (CCDs) were invented in the 1970s and originally found application as memory devices and more recently as shift registers, analog delay lines, particle detectors.

Their light sensitive properties were quickly exploited for imaging applications (TV cameras) and they produced a major revolution in Astronomy. They improved the light gathering power of telescopes by almost two orders of magnitude. Nowadays an amateur astronomer with a CCD camera and a 15 cm telescope can collect as much light as an astronomer of the 1960s equipped with a photographic plate and a 1m telescope.

CCDs work by converting light into a pattern of electronic charge in a silicon chip. This pattern of charge is converted into a video waveform, digitised and stored as an image file on a computer.

The sensitive surface of a CCD is composed of a large number of juxtaposed square or rectangular elements (pixels) sensitive to light; each pixel (a few microns size) creates and accumulates electrical charge proportional to the amount of light it has received and constitutes an elementary point of the image, as does each grain of a photographic film. It is the reading of the accumulated charges in the different pixels that allow the image to be reconstructed.

CCDs used in astronomy are all black and white detectors; it is possible to produce color images by trichromatism (three images of the same object, each one with one of the three fundamental filters and combine these images during analysis). CCDs have practically no detection threshold (unlike films) and the number of electrons generated gives a direct photometric measure.

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• Il rumore intrinseco di un rivelatore e’ tutto ciò che e’

introdotto dal rivelatore stesso e non dal segnale. Determina la soglia inferiore di rilevabilità (con basso rumore intrinseco si osservano sorgenti più deboli)

• Il rumore di un CCD e’ dovuto a:

– Agitazione termica (trascurabile al di sotto di -100°C

– Rumore di lettura del circuito di uscita (ineliminabile)

Se il rivelatore non introducesse rumore, l’unico a restare sarebbe quello della sorgente stessa (RSHOT = √Nfotoni).

• Per fare analisi quantitative il segnale deve superare di almeno 3 volte il rumore

• Il minimo valore di segnale misurabile e’ detto “sensibilità” e si può considerare dato dal

rumore del rivelatore

• “Dynamic range”: the ratio, usually expressed in decibels, of the maximum to the minimum signal that a system can handle (Grey levels). Used to describe the limits of receivers.

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Photoelectric Effect

The effect is fundamental to the operation of a CCD. Atoms in a silicon crystal have electrons arranged in discrete energy bands. The lower energy band is called the Valence Band, the upper band is the Conduction Band. Most of the electrons occupy the Valence band but can be excited into the conduction band by heating or by the absorption of a photon. The energy required for this transition is 1.26 electron volts (eV). Once in this conduction band the electron is free to move about in the lattice of the silicon crystal. It leaves behind a ‘hole’

in the valence band which acts like a positively charged carrier. In the absence of an external electric field the hole and electron will quickly re-combine and be lost. In a CCD an electric field is introduced to sweep these charge carriers apart and prevent recombination.

Increasing energy

Valence Band Conduction Band

1.26eV

Thermally generated electrons are indistinguishable from photo-generated electrons . They constitute a noise source known as ‘Dark Current’ and it is important that CCDs are kept cold to reduce their number.

1.26eV corresponds to the energy of light with a wavelength of 1m. Beyond this wavelength silicon becomes transparent and CCDs constructed from silicon become insensitive.

photon photon

Hole Electron

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CCD Analogy

A common analogy for the operation of a CCD is as follows:

An number of buckets (

Pixels

) are distributed across a field (

Focal Plane of a telescope

)

in a square array. The buckets are placed on top of a series of parallel conveyor belts and collect rain fall (

Photons)

across the field. The conveyor belts are initially stationary, while the rain slowly fills the buckets (

During the course of the exposure

). Once the rain stops (

The camera shutter closes

) the conveyor belts start turning and transfer the buckets of rain , one by one , to a measuring cylinder

(

Electronic Amplifier

) at the corner of the field (

at the corner of the CCD

)

The animation in the following slides demonstrates how the conveyor belts work.

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PHOTONS

PIXELS (BUCKETS

)

CCD COLUMNS

(VERTICAL

CONVEYOR BELTS)

SERIAL REGISTER

(

HORIZONTAL CONVEYOR BELT

)

OUTPUT AMPLIFIER

CCD Analogy

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Exposure finished, buckets now contain samples of rain, i.e. charge has been collected in pixels.

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Conveyor belt starts turning and transfers buckets. Rain collected on the vertical conveyor is tipped into buckets on the horizontal conveyor.

SERIAL REGISTER

(

HORIZONTAL CONVEYOR BELT

)

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Vertical conveyor stops. Horizontal conveyor starts up and tips each bucket in turn into the measuring cylinder .

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`

After each bucket has been measured, the measuring cylinder is emptied , ready for the next bucket load.

OUTPUT

AMPLIFIER

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A new set of empty buckets is set up on the horizontal conveyor and the process is repeated.

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Eventually all the pixels have been measured, the CCD has been read out.

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CCD main functions

Exposure

1. Charge generation: photoelectric effect 2. Charge collection: potential walls (pixels)

Shutter closes

3. Charge transfer: CCD columns shift 4. Readout: serial register and amplification

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Structure of a CCD 1.

The image area of the CCD is positioned at the focal plane of the telescope. An image then builds up that consists of a pattern of electric charge. At the end of the exposure this pattern is then transferred, pixel by pixel, through the serial register to the on-chip amplifier. Electrical connections are made to the outside world via a series of bond pads and thin gold wires positioned around the chip periphery.

Connection pins Gold bond wires Bond pads

Silicon chip Metal, ceramic or plastic package

Image area

Serial register

On-chip amplifier

CCD: 512x512 pixels, 1024x1024, 2048x2048, 4096x4096

Sensitive surface:

•TC211: 2.64x2.64mm²

(192x165 pixels, 13.75x16µm size)

•Thomson 7863: 8.8x6.6mm² (384x288 pixels, 23x23µm)

•SITe SI-003A: 24.6x24.6mm² (1024x1024 pixels, 24x24µm)

•LORAL 2K3eb: 30.7x30.7mm² (2048x2048 pixels, 15x15µm) CCD arrays may be smaller in comparison with standard photography (24x36mm).

“Buttable” CCDs for mosaic.

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Structure of a CCD 2.

CCDs are are manufactured on silicon wafers using the same photo-lithographic techniques used to manufacture computer chips. Scientific CCDs are very big ,only a few can be fitted onto a wafer.

This is one reason that they are so costly.

The photo below shows a silicon wafer with three large CCDs and assorted smaller devices. A CCD has been produced by Philips that fills an entire 6 inch wafer! It is the worlds largest integrated circuit.

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Structure of a CCD 3.

One pixel

Channel stops to define the columns of the image

Transparent

horizontal electrodes to define the pixels vertically. Also used to transfer the charge during readout

Plan View

Cross section

The diagram shows a small section (a few pixels) of the image area of a CCD. This pattern is reapeated.

Electrode

Insulating oxide n-type silicon p-type silicon Every third electrode is connected together. Bus wires running down the edge of the chip make the connection. The channel stops are formed from high concentrations of Boron in the silicon.

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Structure of a CCD 4.

On-chip amplifier at end of the serial register

Cross section of serial register

Image Area

Serial Register

Once again every third electrode in the serial register is connected together.

Below the image area (the area containing the horizontal electrodes) is the ‘Serial register’ . This also

consists of a group of small surface electrodes. There are three electrodes for every column of the image area.

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Structure of a CCD 5.

The serial register is bent twice to move the output amplifier away from the edge

of the chip (Voltage bus). The arrows indicate how charge is transferred through the device.

E dg e of S il ic on

160m

Image Area

Serial Register

Read Out Amplifier

B u s w ir es

Details of a corner of an EEV CCD

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Structure of a CCD 6.

OD

OS RD

R SW

Output Node

Substrate

Output Transistor Reset

Transistor

Summing Well

20m Output Drain (OD)

Output Source (OS) Gate of Output Transistor

Output Node

R

Reset Drain (RD)

Summing Well (SW)

Last few electrodes in Serial Register

Serial Register Electrodes

Details of the on-chip amplifier of a Tektronix CCD and its circuit diagram

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Electric Field in a CCD 1.

The n-type layer contains an excess of electrons that diffuse into the p-layer. The p-layer contains an excess of holes that diffuse into the n-layer. This structure is identical to that of a diode junction.

The diffusion creates a charge imbalance and induces an internal electric field. The electric potential reaches a maximum just inside the n-layer, and it is here that any photo-generated electrons will collect.

All science CCDs have this junction structure, known as a ‘Buried Channel’. It has the advantage of

keeping the photo-electrons confined away from the surface of the CCD where they could become trapped.

It also reduces the amount of thermally generated noise (dark current).

n p

Electric potential

Potential along this line shown in graph above.

Electric potential

Cross section through the thickness of the CCD

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Electric Field in a CCD 2.

During integration of the image, one of the electrodes in each pixel is held at a positive potential. This further increases the potential in the silicon below that electrode and it is here that the photoelectrons are accumulated. The neighboring electrodes, with their lower potentials, act as potential barriers that define the vertical boundaries of the pixel. The horizontal boundaries are defined by the channel stops.

n p

Electric potential

Region of maximum potential

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pixel boundary

Charge packet

p-type silicon n-type silicon

SiO2 Insulating layer Electrode Structure

pixel boundary

incoming photons

Charge Collection in a CCD.

Photons entering the CCD create electron-hole pairs. The electrons are then attracted towards the most positive potential in the device where they create ‘charge packets’. Each packet corresponds to one pixel

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Charge Transfer in a CCD 1.

In the following few slides, the implementation of the ‘conveyor belts’ as actual electronic structures is explained.

The charge is moved along these conveyor belts by modulating the voltages on the electrodes positioned on the surface of the CCD. In the following illustrations, red electrodes

are held at a positive potential, black ones are held at a negative potential.

1

2

3

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1 2 3

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1 2 3

+5V 0V -5V +5V 0V -5V +5V 0V -5V

1 2

3 Charge Transfer in a CCD 8.

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On-Chip Amplifier 1.

OD

OS RD

R SW

Output

Node Output

Transistor Reset

Transistor

Summing Well

+5V 0V -5V +10V

0V

R SW

--end of serial register

V

_out

V

_out

The on-chip amplifier measures each charge packet as it pops out the end of the serial register.

The measurement process begins with a reset of the ‘reset node’. This removes the charge remaining from the previous pixel. The reset node is in fact a tiny capacitance (< 0.1pF) RD and OD are held at

constant voltages

(The graphs above show the signal waveforms)

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On-Chip Amplifier 2.

OD

OS RD

R SW

Output

Node Output

Transistor Reset

Transistor

Summing Well

+5V 0V -5V +10V

0V

R SW

V

_out

V

_out

The charge is then transferred onto the Summing Well.

V

_out is now at the ‘Reference level’

There is now a few tens of µs wait while external circuitry measures this ‘reference’ level.

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On-Chip Amplifier 3.

OD

OS RD

R SW

Output

Node Output

Transistor Reset

Transistor

Summing Well

+5V 0V -5V +10V

0V

R SW

V

_out

V

_out

This action is known as the ‘charge dump’

The voltage step in

V

_outis as much as several V for each electron contained in the charge packet.

The charge is then transferred onto the output node.

V

_out now steps down to the ‘Signal level’

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On-Chip Amplifier 4.

OD

OS RD

R SW

Output

Node Output

Transistor Reset

Transistor

Summing Well

+5V 0V -5V +10V

0V

R SW

V

_out

V

_out

V

_out is now sampled by external circuitry for up to a few tens of microseconds.

The sample level - reference level will be proportional to the size of the input charge packet.

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Use of CCD Cameras

Practical considerations of using and building CCD cameras

Nik Szymanek

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Spectral Sensitivity of CCDs

The graph below shows the transmission of the atmosphere when looking at objects at the zenith. The atmosphere absorbs strongly below about 330nm, in the near ultraviolet part of the spectrum. An ideal CCD should have a good sensitivity from 330nm to approximately 1000nm, at which point silicon, from which CCDs are manufactured, becomes transparent and therefore insensitive.

Wavelength (Nanometers)

Transmission of Atmosphere

Over the last 25 years of development, the sensitivity of CCDs has improved enormously, to the point where almost all of the incident photons across the visible spectrum are detected.

CCD sensitivity has been improved using two main techniques: ‘thinning’ and the use of anti-reflection coatings. They will be now explained in more details.

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A CCD sensitivity (Quantum Efficiency, QE) is expressed as the number of electrons

produced per incident photon.

To improve UV sensitivity:

•Anti-reflection coating:

fluorescent coat excited by UV light, re-emits in the green

•Thinning: Back-lit CCDs Pixel capacity:10⁵-10⁶ electrons. Saturation,

“blooming” (bright stars)

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Thick Front-side Illuminated CCD

These are cheap to produce using conventional wafer fabrication techniques. They are used in consumer imaging applications. Even though not all the photons are detected, these devices are still more sensitive than photographic film.

They have a low Quantum Efficiency due to the reflection and absorption of light in the surface electrodes. Very poor blue response. The electrode structure prevents the use of an Anti-reflective coating that would otherwise boost performance.

The amateur astronomer on a limited budget might consider using thick CCDs. For

professional observatories, the economies of running a large facility demand that the detectors be as sensitive as possible; thick front-side illuminated chips are seldom if ever used.

n-type silicon p-type silicon

Silicon dioxide insulating layer Polysilicon electrodes

Incoming photons

625m

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Anti-Reflection Coatings 1

n

of air or vacuum is 1.0, glass is 1.46, water is 1.33, Silicon is 3.6. Using the above equation we can show that window glass in air reflects 3.5% and silicon in air reflects 32%. Unless we take steps to eliminate this reflected portion, then a silicon CCD will at best only detect 2 out of every 3 photons.

n

s

n

t

2 =

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Anti-Reflection Coatings 3

The graph below shows the reflectivity of an EEV 42-80 CCD. These thinned CCDs were designed for a maximum blue response and it has an anti-reflective coating optimised to work at 400nm. At this wavelength the reflectivity falls to approximately 1%.

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Thinned Back-side Illuminated CCD

The silicon is chemically etched and polished down to a thickness of about 15microns. Light enters from the rear and so the electrodes do not obstruct the photons. The QE can approach 100% .

These are very expensive to produce since the thinning is a non-standard process that reduces the chip yield. These thinned CCDs become transparent to near infra-red light and the red response is poor. Response can be boosted by the application of an anti-reflective coating on the thinned rear-side. These coatings do not work so well for thick CCDs due to the surface bumps created by the surface electrodes.

Almost all Astronomical CCDs are Thinned and Backside Illuminated.

n-type silicon p-type silicon

Silicon dioxide insulating layer Polysilicon electrodes

Incoming photons

Anti-reflective (AR) coating

15m

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Quantum Efficiency Comparison

The graph below compares the quantum of efficiency of a thick frontside illuminated CCD and a thin backside illuminated CCD.

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‘Internal’ Quantum Efficiency

If we take into account the reflectivity losses at the surface of a CCD we can produce a graph showing the ‘internal QE’ : the fraction of the photons that enter the CCDs bulk that actually produce a

detected photo-electron. This fraction is remarkably high for a thinned CCD. For the EEV 42-80 CCD, shown below, it is greater than 85% across the full visible spectrum. Todays CCDs are very close to being ideal visible light detectors!

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Appearance of CCDs

The fine surface electrode structure of a thick CCD is clearly visible as a multi-coloured interference pattern. Thinned Backside Illuminated CCDs have a much planer surface appearance. The other notable distinction is the two-fold (at least) price difference.

Kodak Kaf1401 Thick CCD MIT/LL CC1D20 Thinned CCD

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Computer Requirements 1.

Computers are required firstly to coordinate the sequence of clock signals that need to be sent to a CCD and its signal processing electronics during the readout phase, but also for data collection and the

subsequent processing of the images.

The CCD Controller

In this first application, the computer is an embedded system running in a ‘CCD controller’. This controller will typically contain a low noise analogue section for amplification and filtering of the CCD video waveform, an analogue to digital converter, a high speed processor for clock waveform generation and a fibre optic transceiver for receipt of commands and transmission of pixel data.

An astronomical system might require clock signals to be generated with time resolutions of a few tens of nanoseconds. This is typically done using Digital Signal Processing (DSP) chips running at 50Mhz. Clock sequences are generated in software and output from the DSP by way of on-chip parallel ports. The most basic CCD design requires a minimum of 7 clock signals. Perhaps 5 more are required to coordinate the operation of the signal processing electronics. DSPs also contain several on-chip serial ports

which can be used to transmit pixel data at very high rates. DSPs come with a small on-chip memory for the storage of waveform generation tables and software. Less time critical code , such as routines to initialise the camera and interpret commands can be stored in a few KB of external RAM. The computer running in the CCD controller is thus fast and of relatively simple design. A poorly performing processor here could result in slow read out times and poor use of telescope resources. Remember that when a CCD is reading out the telescope shutter is closed and no observations are possible. For an amateur observer using a small CCD with a fast readout time, a slow CCD controller may not be such a disadvantage; there are not so many pixels to process.

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Computer Requirements 2.

The Data Acquisition System(DAS)

This will be typically based around a SUN SPARC workstation which is a high-end desktop computer. Pixel data will be received from the CCD controller by way of a fibre optic. The hardware in such a system will be cheap and ‘off-the shelf’, the only speciality item being the high speed fibre optic transceiver card.

The hardware may typically consist of a Sparc Ultra 6 workstation, 500Mb of RAM, a 9GB hard-drive and a DAT drive. There will also be a high speed Ethernet card for connection to the observatory Local Area Network.The software required to carry out the data acquisition task is typically developed in-house by each observatory and represents the major cost of such a system. It will provide an easy-to-use interface (typically graphics based) between observer and instrument. Its complexity will be further increased by the need to talk to other telescope systems such as the Telescope Control System. This will allow information on the pointing of the telescope to be stored alongside the pixel data as a ‘file header’.

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Computer Requirements 3.

Image Processing Computers

These are used for reduction and analysis of the astronomical data. Many astronomers process their data in real-time, i.e. they may be analysing one exposure whilst the next exposure is actually been taken. Others will take a cursory look at their data in real time but leave the heavy image processing tasks for when they return to their home institution. With large mosaic cameras producing very large data files, a high end system is required.

A typical system would be :

A PC with a 1GHz CPU

Enough RAM for at least 2 images , using 4 bytes per pixel (for a mosaic camera this could run to 500MBytes)

At least 100 GBytes (300GBytes would be better) of local hard disc space

If we use such a system to analyse images from a four chip CCD mosaic containing 36 Million pixels, the following performance would be obtained :

Linearisation, bias subtraction and flat-fielding : ~150 sec de-fringing : ~300 sec object detection and star/galaxy separation : ~300 sec

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Computer Requirements 4.

Image Processing Computers (Contd.)

This professional system is unusual in its high demands on disc space and RAM. The processor speed , however, is the same as that found in current PCs costing a few thousand dollars. An amateur observer with a small 1K square CCD camera will find a medium level PC quite sufficient for operation of the camera and for image processing.

The system specs would typically be:

Pentium III 500 MHz processor,

256 MB RAM

32 MB video memory

30GB Hard Drive

CD Writer

a 19” monitor (twin monitors are even better, one for images , one for text)

For operation of the camera the bottleneck is often the data transfer between camera and PC. For image processing applications such as Maximum Entropy or Lucy-Richardson de-convolution

(a form of image sharpening), a high speed PC is needed.

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Schematic view of CCD camera electronics

Clocks

CCD

Amplifier

Digital converter

Cooling system

I/O

interface Computer

CCD camera Acquisition

commands

Image

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Image defects in CCDs

Blooming: limited capacity of a CCD pixel

Dark columns: traps

Cosmic rays

Bright columns: traps

Hot Spots: non uniformity of structure

Electronic processing

Calibration exposures

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pixel boundary

Photons

Blooming 1.

The charge capacity of a CCD pixel is limited: when a pixel is full, the charge starts to leak into adjacent pixels. This process is known as ‘Blooming’.

Photons

Overflowing charge packet

Spillage Spillage

pixel boundary

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Blooming 2.

The diagram shows one column of a CCD with an over-exposed stellar image focused on one pixel.

The channel stops shown in yellow prevent the charge spreading sideways. The charge confinement provided by the electrodes is lower so the charge spreads vertically up and down a column.

The capacity of a CCD pixel is known as the ‘Full Well’. It is dependent on the physical area of the pixel. For Tektronix

CCDs, with pixels measuring 24m x 24m it can be as much as 300,000 electrons. Bloomed images will be seen particularly on nights of good seeing where stellar images are more compact . In reality, blooming is not a big problem for professional

astronomy. For those interested in pictorial work, however, it can be a nuisance.

Flow of bloomed charge

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Blooming 3.

Bloomed star images The image below shows an extended source with bright embedded stars. Due to the long exposure required to bring out the nebulosity, the stellar images are highly overexposed and create bloomed images.

(The image is from a CCD mosaic and the black strip down the center is the space between adjacent detectors)

M42

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Dark columns

Unless one pays a huge amount it is generally difficult to obtain a CCD free of image defects.

The first kind of defect is a ‘dark column’. Their locations are identified from flat field exposures.

Flat field exposure of an EEV42-80 CCD

Dark columns are caused by ‘traps’ that block the vertical transfer of charge during image readout. The CCD shown at left has at least 7 dark columns, some grouped together in adjacent clusters.

Traps can be caused by crystal boundaries in the silicon of the CCD or by manufacturing defects.

Although they spoil the chip cosmetically, dark columns are not a big problem for astronomers. This chip has 2048 image columns so 7 bad columns represents a tiny loss of data.

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Cosmic rays, Hot spots, Bright columns

Bright columns are also caused by traps . Electrons contained in such traps can leak out during readout causing a vertical streak Hot Spots are pixels with higher than normal dark current. Their brightness increases linearly with exposure times

Cosmic rays are unavoidable. Charged particles from space or from radioactive traces in the material of the camera can cause ionisation in the silicon. The electrons produced are indistinguishable from photo-generated electrons.

Approximately 1-2 cosmic rays/cm²/minute will be seen.

A typical event will be spread over a few adjacent pixels and contain several thousand electrons.

Somewhat rarer are light-emitting defects which are hot spots that act as tiny LEDS and cause a halo of light on the chip.

There are three more image defect types : cosmic rays, bright columns and hot spots.

Their locations are shown in the image below, which is a lengthy exposure taken in the dark (a ‘Dark Frame’)

Cosmic rays Cluster of Hot Spots Bright Column

900s dark exposure of an EEV42-80 CCD

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Electronic processing

M51

Some defects can arise from the read-out processing electronics. This negative image has a bright line in the first image row.

Dark column

Hot spots and bright columns

Bright first image row caused by incorrect operation of signal processing electronics.

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Calibration exposures

Bias frames

CCD defects, readout noise

exposure time = 0

Flat field

pixel non-uniformity, dust, …

exposure time ≈ medium (half saturation)

Dark frames

dark current, defects

exposure time = science image

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Biases, Flat Fields and Dark Frames 1.

These are three types of calibration exposures that must be taken with a scientific CCD camera, generally before and after each observing session. They are stored alongside the science images

and combined with them during image processing. These calibration exposures allow us to compensate for certain imperfections in the CCD. As much care needs to be exercised in obtaining these images as for the actual scientific exposures. Applying low quality flat fields and bias frames to scientific data can degrade its quality, rather than improve it.

Bias Frames (CCD defects, readout noise)

A bias frame is an exposure of zero duration taken with the camera shutter closed. It represents the zero point or base-line signal from the CCD. Rather than being completely flat and featureless, the bias frame may contain some structure. Any bright image defects in the CCD will of course show up, there may also be slight gradients in the image caused by limitations in the signal processing electronics of the camera.

It is normal to take about 5 bias frames before a night’s observing. These are then combined using an image processing algorithm that averages the images, pixel by pixel, rejecting any pixel values that are appreciably different from the other 4. This can happen if a pixel in one bias frame is affected by a cosmic ray event. It is unlikely that the same pixel in the other 4 frames would be similarly affected so the resultant ‘master bias’, should be uncontaminated by cosmic rays. Taking a number of biases and then averaging them also reduces the amount of noise in the bias images. Averaging 5 frames will reduce the amount of read noise (electronic noise from the CCD amplifier) in the image by the square-root of 5.

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Biases, Flat Fields and Dark Frames 2.

Flat Fields (pixel non-uniformity, dust, etc.)

Some pixels in a CCD will be more sensitive than others (non-uniformity due to slightly different pixel dimensions and depth because of manufacturing process). In addition there may be dust spots on the surface of either the chip, the window of the camera or the colored filters mounted in front of the camera. A star focused onto one part of a chip may therefore produce a lower signal than it might do elsewhere. These variations in sensitivity across the surface of the CCD must be calibrated out or they will add noise to the image. The way to do this is to take a ‘flat-field ‘ image: an image in which the CCD is evenly illuminated with light. Dividing the science image, pixel by pixel, by a flat field image will remove these non-uniformities and sensitivity variations very effectively.

Since some of these variations are caused by shadowing from dust spots, it is important that the flat fields are taken shortly before or after the science exposures as the dust may move around. As with biases, it is normal to take several flat field frames and average them to produce a ‘Master’.

A flat field is taken by pointing the telescope at an extended, evenly illuminated source. The twilight sky or the inside of the telescope dome are the usual choices. An exposure time is chosen that gives pixel values about halfway to their saturation level, i.e. a medium level exposure.

Dark Frames (dark current, defects)

Dark current is generally absent from professional cameras since they are operated at low temperature, using liquid nitrogen as a coolant. Amateur systems running at higher temperatures will have some dark current and its effect must be minimized by obtaining ‘dark frames’ at the beginning of the observing run. These are exposures with the same duration as the science frames but taken with the camera shutter closed. These are later subtracted from the science frames. Again, it is normal to take several dark frames and combine them to form a Master,

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Dark Frame Flat Field

A dark frame and a flat field from the same EEV42-80 CCD are shown below. The dark frame shows a number of bright defects on the chip. The flat field shows a cross patterning on the chip

created during manufacture and a slight loss of sensitivity in two corners of the image. Some dust spots are also visible.

Biases, Flat Fields and Dark Frames 3.

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Flat Field Frame

Bias Frame

Flat -Bias Science -Dark

Output Image

Flat-Bias Science -Dark

Dark Frame Science Raw Image

Biases, Flat Fields and Dark Frames 4.

All raw images contain undesirable effects, inherent to the CCD system technology used; preprocessing consists in “cleaning” the raw image of these effects.

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Bias Frames and Flat Fields 5.

Flat Field Image Bias Image

Flat -Bias Science -Bias

Output Image

Flat-Bias Science -Bias

Science Frame

In the absence of dark current, the process is slightly simpler :

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Pixel Size and Binning 1.

It is important to match the size of a CCD pixel to the focal length of the telescope. Atmospheric seeing places a limit on the sharpness of an astronomical image for telescope apertures above 15cm. Below this aperture, the images will be limited by diffraction effects in the optics. In excellent seeing conditions, a large telescope can produce stellar images with a diameter of 0.6 arc-seconds. In order to record all the information present in such an image, two pixels must fit across the stellar image; the pixels must subtend at most

0.3 arc-seconds on the sky. This is the ‘Nyquist criteria’. If the pixels are larger than 0.3 arc-seconds the Nyquist criteria is not met, the image is under-sampled and information is lost. The Nyquist criteria also applies to the digitisation of audio waveforms. The audio bandwidth extends up to 20KHz , so the Analogue to Digital Conversion rate needs to exceed 40KHz for full reproduction of the waveform. Exceeding the Nyquist criteria leads to ‘over-sampling’.This has the disadvantage of wasting silicon area; with improved matching of detector and optics a larger area of sky could be imaged.

Under-sampling an image can produce some interesting effects. One of these is the introduction of features that are not actually present. This is occasionally seen in TV broadcasts when, for example, the fine-patterned shirt of an interviewee breaks up into psychedelic bands and ripples. In this example, the TV camera pixels are too big to record the fine detail present in the shirt.

This effect is known as ‘aliasing’.

Nyquist/Shannon Theorem: f

_c

≥2f

_max

(f

_c

=sampling freq, f

_max

=signal max freq)

Nyquist Sampling

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Aliasing

When a sinusoidal signal is measured or sampled at regular but not sufficiently close intervals, one will obtain the same sequence of samples that would be obtained from a sinusoid of a lower frequency. Specifically, if a sinusoid of frequency f (in cycles per second for a time-varying signal, or in cycles per centimeter for space- varying signal) is sampled f_s samples per second or per centimeter, the resulting samples will also be compatible with a sinusoid of frequency (Nf_s-f) and one of frequency (Nf_s+f), for any integer N. If f_s>2f , the lowest of these image frequencies will be the original signal frequency, but otherwise it will not. In the case that

f_s<2f<2f_s, the lowest image frequencies will be at (f_s-f) , the lowest image frequency in a sense masquerades as the sinusoid that was sampled and is called an alias of the sinusoid that was actually sampled, albeit

inadequately sampled.

One way to avoid such aliasing is to make sure that the signal does not contain any sinusoidal component with a frequency equal to or greater than f_s/2. This condition is sometimes called the Nyquist criterion, and is

equivalent to saying that the sampling frequency f_s must be high enough; either greater than twice the highest frequency or some other more complicated criterion.

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Aperture in mm X f-number

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Pixel Size and Binning 2.

Example 1.

The William Herschel Telescope, with a 4.2m diameter primary mirror and a focal ratio of 3 is to be used for prime focus imaging. What is the optimum pixel size assuming that the best seeing at the telescope site is 0.7 arc-seconds ?

First we calculate the ‘plate-scale’ in arc-seconds per millimeter at the focal plane of the telescope.

Plate Scale (arc-seconds per mm) = =16.4 arc-sec per mm

(Here the factor 206265 is the number of arc-seconds in a Radian)

Next we calculate the linear size at the telescope focal plane of a stellar image (in best seeing conditions) Linear size of stellar image = 0.7 / Plate Scale = 0.7/ 16.4 = 43 microns.

To satisfy the Nyquist criteria, the maximum pixel size is therefore 21microns. In practice, the nearest pixel size available is 13.5 microns which leads to a small degree of over-sampling.

Matching the Pixels to the telescope

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Pixel Size and Binning 3.

Example 2.

An Amateur telescope with a 20cm aperture and a focal ratio of 10 is to be used for imaging. The best seeing conditions at the observing site will be 1 arc-second. What is the largest pixel size that can be used?

Plate Scale (arc-seconds per mm) = =103 arc-sec per mm

Linear size of stellar image = 1 / Plate Scale = 1/ 103 = 9.7 microns.

To satisfy the Nyquist criteria, the maximum pixel size is therefore 5 microns. This is about the lower limit of available pixel sizes.

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Pixel Size and Binning 4.

In the first example we showed that with 13.5µm pixels the system exceeded the Nyquist Criteria even on nights with exceptionally good sub-arcsecond seeing. If we now suppose that the seeing is 2

arc-seconds, the size of a stellar image will increase to 120µm on the detector. The image will now be grossly over-sampled. (One way to think of this is that the image is less sharp and therefore requires fewer pixels to record it). It would be more efficient now for the astronomer to switch to a detector with larger pixels since the resultant image files would be smaller, quicker to read out and would occupy less disk space.

There is a way to read out a CCD so as to increase the effective pixel size, this is known as ‘Binning’. With binning we can increase pixel size arbitrarily. In the limit we could even read out the CCD as a single large pixel. Astronomers will more commonly use 2 x 2 binning which means that the charge in each 2 x 2 square of adjacent pixels is summed on the chip prior to delivery to the output amplifier. One important advantage of ‘on-chip binning’ is that it is a noise free process.

Binning is done in two distinct stages : vertical binning and horizontal binning. Each may be done without the other to yield rectangular pixels.

Binning

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Pixel Size and Binning 5.

This is done by summing the charge in consecutive rows .The summing is done in the serial register. In the case of 2 x 2 binning, two image rows will be clocked consecutively into the serial register prior to the serial register being read out. We now go back to the conveyor belt analogy of a CCD. In the following animation we see the bottom two image rows being binned.

Stage 1 :Vertical Binning

Charge packets

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Pixel Size and Binning 6.

The first row is transferred into the serial register

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Pixel Size and Binning 7.

The serial register is kept stationary ready for the next row to be transferred.

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Pixel Size and Binning 8.

The second row is now transferred into the serial register.

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Pixel Size and Binning 9.

Each pixel in the serial register now contains the charge from two pixels in the image area. It is thus important that the serial register pixels have a higher charge capacity. This is achieved by giving them a larger physical size.

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Pixel Size and Binning 10.

This is done by combining charge from consecutive pixels in the serial register on a special electrode positioned between serial register and the readout amplifier called the Summing Well (SW).

The animation below shows the last two pixels in the serial register being binned : Stage 2 :Horizontal Binning

1 2 3

SW

Output Node

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Pixel Size and Binning 11.

Pixel Size and Binning 14.

The charge from the first pixel is now stored on the summing well.

1 2 3

SW

Output Node

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Pixel Size and Binning 15.

The serial register continues clocking.

1 2 3

SW

Output Node

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Pixel Size and Binning 16.

1 2 3

SW

Output Node

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Pixel Size and Binning 17.

The SW potential is set slightly higher than the serial register electrodes.

1 2 3

SW

Output Node

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Pixel Size and Binning 18.

1 2 3

SW

Output Node

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Pixel Size and Binning 19.

The charge from the second pixel is now transferred onto the SW. The binning is now complete and the combined charge packet can now be dumped onto the output node (by pulsing the voltage on SW low for a microsecond) for measurement.

Horizontal binning can also be done directly onto the output node if a SW is not present but this can increase the read noise.

1 2 3

SW

Output Node

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Pixel Size and Binning 20.

Finally the charge is dumped onto the output node for measurement

1 2 3

SW

Output Node

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Advanced CCD Techniques

CCD Imaging

Nik Szymanek

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Integrating and Video CCD Cameras.

There is a difference in the geometry of an Integrating CCD camera compared

to a Video CCD camera. An integrating camera, such as is used for most astronomical applications, is designed to stare at an object over an exposure time of many minutes.

When readout starts and the charge is transferred out of the image area, line by line, into the serial register, the image area remains light sensitive. Since the readout can take as long as a minute, if there is no shutter, each stellar image will be drawn out into a line. An external shutter is thus essential to prevent smearing. These kind of CCDs are known as ‘Slow Scan’.

A video CCD camera is required to read out much more rapidly. A video CCD may be used by the astronomers as a finder-scope to locate objects of interest and ensure that the telescope is actually pointed at the target or it may be used for auto-guiding. These cameras must read out much more quickly, perhaps several times a second. A mechanical shutter

operating at such frame rates could be unreliable. The geometry of a video CCD, however , incorporates a kind of electronic shutter on the CCD and no external shutter is required.

These kind of CCDs are known as ‘Frame Transfer’.

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Slow Scan CCDs 1.

The most basic geometry of a Slow-Scan CCD is shown below. Three clock lines

control the three phases of electrodes in the image area, another three control those in the serial register. A single amplifier is located at the end of the serial register. The full image area is available for imaging. Because all the pixels are read through a single output, the readout speed is relatively low. The red line shows the flow of charge out of the CCD.

Image area clocks

Serial Register clocks Serial Register Output Amplifier

Image Area

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Slow Scan CCDs 2.

A slightly more complex design uses 2 serial registers and 4 output amplifiers. Extra clock lines are required to divide the image area into an upper and lower section. Further clock lines allow independent operation of each half of each serial register. It is thus possible to read out the image in four quadrants simultaneously, reducing the readout speed by a factor of four.

Upper Image area clocks

Lower Image area clocks

Amplifier C

Amplifier A Amplifier B

Amplifier D

Serial clocks C Serial clocks D Serial clocks A Serial clocks B

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Slow Scan CCDs 3.

There are certain drawbacks to using this ‘split-frame readout’ method. The first is that each amplifier will have slightly different characteristics. It may have a slightly different gain or a differing linearity. Reconstructing a single image from the four sub-images can be an image processing nightmare and unless the application demands very high readout speed, most astronomers would wait slightly longer for an image read out through a single amplifier.

Another drawback is cost. CCDs that have all of their output amplifiers working are rare and come at a premium price.

Most CCDs are designed with multiple outputs. Even if only one of the working outputs is actually used, the others provide valuable backups should there be, for any reason, an amplifier failure.

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Video CCDs 1.

Image area clocks

Store area clocks

Amplifier

Serial clocks Image area

Store area

In the split frame CCD geometry, the charge in each half of the image area could be shifted independently. Now imagine that the lower image area is covered with an opaque mask. This mask could be a layer of aluminum deposited on the CCD surface or it could be an external

mask. This geometry is the basis of the ‘Frame transfer’ CCD that is used for high frame rate video applications. The area available for imaging is reduced by a half. The lower part of the image becomes the ‘Store area’.

Opaque mask

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Video CCDs 2.

The operation of a Split Frame Video CCD begins with the integration of the image in the image area.

Once the exposure is complete the charge in the image area is shifted down into the store area

beneath the light proof mask. This shift is rapid; of the order of a few milliseconds for a large CCD.

The amount of image smear that will occur in this time is minimal (remember there is no external shutter).

Integrating Galaxy Image

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Video CCDs 3.

Once the image is safely stored under the mask, it can then be read out at leisure. Since we can independently control the clock phases in the image and store areas, the next image can be

integrated in the image area during the readout. The image area can be kept continuously integrating and the detector has only a tiny ‘dead time’ during the image shift. No external shutter is required but the effective size of the CCD is cut by a half.

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Correlated Double Sampler (CDS) 1.

The video waveform output by a CCD is at a fairly low level : every photo-electron in a pixel charge packet will produce a few micro-volts of signal. Additionally, the waveform is

complex and precise timing is required to make sure that the correct parts are amplified and measured.

The CCD video waveform , as introduced in Activity 1, is shown below for the period of one pixel measurement

V

_out

t

Reset feedthrough Reference level Charge dump Signal level

The video processor must measure , without introducing any additional noise, the Reference level and the Signal level. The first is then subtracted from the second to yield the output signal voltage proportional to the number of photo-electrons in the pixel under measurement. The best way to perform this processing is to use a ‘Correlated Double Sampler’ or CDS.

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OD

OS RD

R

. ^-1 ^ADC

Pre-Amplifier

CCD On-chip Amplifier

Inverting Amplifier

Integrator Reset switch

Input Switch Polarity Switch

Correlated Double Sampler (CDS) 2.

The CDS design is shown schematically below. The CDS processes the video waveform and outputs a digital number proportional to the size of the charge packet contained in the pixel being read. There should only be a short cable length between CCD and CDS to minimise noise.The CDS minimises the read noise of the CCD by eliminating ‘reset noise’. The CDS contains a high speed analogue processor containing computer controlled switches. Its output feeds into an Analogue to Digital Converter (ADC).

C om pu te r B us

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t

₀

t

₀

Correlated Double Sampler (CDS) 3.

Output wave-form of CCD

The CDS starts work once the pixel charge packet is in the CCD summing well and the CCD reset pulse has just finished. At point t₀the CCD wave-form is still affected by the reset pulse and so the CDS remains disconnected from the CCD to prevent this disturbing the video processor.

-1

Output voltage of CDS

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t

₁

t

₂

t

₁

t

₂

-1

Correlated Double Sampler (CDS) 4.

Between t₁and t₂the CDS is connected and the ‘Reference ‘ part of the waveform is sampled.

Simultaneously the integrator reset switch is opened and the output starts to ramp down linearly.

Reference window

Correlated Double Sampler (CDS) 5.

Between t₂and t₃the ‘charge dump’ occurs in the CCD. The CCD output steps negatively by an amount proportional to the charge contained in the pixel. During this time the CDS is disconnected.

₄

-1

Correlated Double Sampler (CDS) 6.

Between t₃and t₄the CDS is reconnected and the ‘signal’ part of the wave-form is sampled. The input to the integrator is also ‘polarity switched’ so that the CDS output starts to ramp-up linearly. The width of the signal and sample windows must be the same. For Scientific CCDs this can be anything between 1 and 20 microseconds. Longer widths generally give lower noise but of course increase the read-out time.

Signal window

The CDS is then once again disconnected and its output digitised by the ADC. This number , typically a 16 bit number (with a value between 0 and 65535) is then stored in the computer memory. The CDS then starts the whole process again on the next pixel. The integrator output is first zeroed by closing the reset switch. To process each pixel can take between a fraction of a microsecond for a

TV rate CCD and several tens of microseconds for a low noise scientific CCD.

The type of CDS is called a ‘dual slope integrator’.

A simpler type of CDS known as a ‘clamp and sample’

only samples the waveform once for each pixel.

It works well at higher pixel rates but is noisier

than the dual slope integrator at lower pixel rates. ^V

oltage to be digitised

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Noise Sources in a CCD Image 1.

Noise level sets the smallest detectable signal (faint sources) so it must be properly controlled.

The main noise sources found in a CCD are :

1. READ OUT NOISE

Caused by electronic noise in the CCD output transistor and possibly also in the external circuitry.

Read noise places a fundamental limit on the performance of a CCD. It can be reduced at the expense of increased read out time. Scientific CCDs have a read out noise of 2-3 electrons RMS.

2. DARK CURRENT

Caused by thermally generated electrons in the CCD even in the absence of lighting, increasing with temperature and integration time.

Consequences: generates “thermal noise” limiting the detection of faint sources and can even saturate the pixels. Eliminated by cooling down the CCD usually at ~-100ºC.

3. PHOTON NOISE

Also called ‘Shot Noise’. It is due to the fact that the CCD detects photons. Photons arrive in an unpredictable fashion described by Poisson statistics. This unpredictability causes noise.

4. PIXEL RESPONSE NON-UNIFORMITY

Defects in the silicon and small manufacturing defects can cause some pixels to have a higher sensitivity than their neighbours. This noise source can be removed by ‘Flat Fielding’.

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