The evolution of detectors: from photography to CCD s

Transcription

1 The evolution of detectors: from photography to CCD s Michel Dennefeld Institut d Astrophysique de Paris and University P. et M. Curie (Paris 6) and thanks to Henry Mc Cracken...

2 Part I : History and progress... The human eye: -2.5 cm diameter lens -130 millions pixels -125 M cones (day-time, tri-colour vision) ~ 140 FOV, 1 resolution -5 M rods (night-time, grey vision) -Spectral response: -Daytime: peak at ~ 550 nm -Nightime: slightly bluer

3 Transparency and energy versus λ hν~ GeV-keV ~ 1 ev? I=[a(sinωt-φ)]²

4 Modern detectors The principle is the same for all detectors: transform a photon into a (photo)-electron! What makes the difference, is what happens then to this electron... hν = E extr + ½ mv 2 Wavelength threshold: λ (μ) < 1.24/ E (ev) e.g. Silicon: 1.1 μ; Photo: ~ 0.45 μ -If the electron «escapes» the material, you have photoelectric systems Requires additional extraction energy (thus less red sensitivity), but you can then act on the electron (amplification) -If the electron remains inside, you have photoconductors Higher efficiency, but more difficult to handle

5 Photographic plates AgBr crystals: Br - + hv gives Br + e - and Ag + + e - gives (unstable) Ag The development enhances the quantity of Ag where it exists. Problems: very non-linear process; impurities and traps compete for the electrons, hence very low efficiency. The latter can be improved with sensitising processes, but the DQE hardly comes over 3%! It was the first objective and permanent record of astronomical phenomena (~ ) Long exposure times are possible thus faint objects can be detected Exists in large format: many sky surveys carried out, as recently as the 1990s. e.g. the Palomar Sky Surveys (DSS...) in B (103aO) and R (IIaF); in the South, IIIaJ, etc... But what are QE, DQE, 103aO, etc...??

6 Fondamental parameters for a detector... QE (noted η) = <number of electrons/photons>, is << 1! DQE = (input S/N) 2 / (output S/N) 2 thus DQE is η (in Poisson statistics, the S/N is equal to n ) Sensitivity: NEP, input power needed to provide a S/N of 1 Its dependance with wavelength... Linearity, dynamical range, MTF

7 Parameters of photographic plates How to improve sensitivity? -Sensitizing (= reduce traps) - Higher energy electrons (photocathodes + high voltage = electronographic systems)

8 Photoelectric detectors First application: Photomultiplier tubes ( s) = single pixel detectors! High-voltage (~ 20 kv) acceleration and amplification of electrons. Initially measured output currents, later used as photon-counting systems. More sensitive than photography (~20% quantum efficiency) Extremely stable, linear, broad band devices which made possible the first highly accurate photometric systems Flux calibration of photomultiplier based systems is still more accurate than modern array imaging detectors!

9 How to get 2D images with photoelectric systems? Electronographic cameras: Electrostatic, or magnetic focusing, +photographic plate output But photocathodes need to remain in vacuum... (Lallemand, Kron, Spectracon cameras...) TV tubes, but not very stable, and poor PSF Combined systems lead to Photon-Counting devices.

10 How do Photon Counting systems work? High e- amplification: bright patch of light at the output Centroiding of the light patch, at a fraction of (output) pixel. Coordinates of the event are stored in the memory: independant of the intensity of the patch! Thus each event has the same weight Single photo-electron is detected if amplification high enough No noise, but saturation limit Deadtime between events (e.g. decay time of output phosphor)

11 Modern photon-counting systems Micro-channel plates: more compact, and lighter than photo-tubes. Various readout modes...the best is: MAMA (Multi-Anode Multichannel Arrays) Photon counting systems: -Pro s: - Single photon detected - Equal weight for each event - No noise -Con s: - Sensitive to high-flux (deadly!!) - Low QE (limited by photocathode, ~20%) - DQE ~ 7% (e.g. HST/STIS)

12 How to improve the Quantum Efficiency? and its wavelength dependance? ~20% QE limit for photo-emissivity, and poor red-sensitivity Semi-conductors have better QE, but electrons trapped inside One thus needs to find ways to «read» the electrons...

13 How a semi-conductor works... Exemple of a photo-diode - Energy diagram -Read-out mode Many capacitors in parallel, thus high read-out noise (NEC ~C v /q [4kTBR 1 ] )

14 How to do it on a 2D image: CCD Readout Thomson 2048 X 2048 CCD Digitize via Analog to digital (A/D) converter Horizontal register = Serial register 14

15 Charge Coupled Devices (CCD s) To reduce C v : no external connections, and no physical limit between pixels! Only one connection at the extremity. The free photo-electrons in the conducting band are then shuffled to the output electrodes by applying electrical pulses. The output signal is amplified and converted from electron units to digital units with a conversion factor or Gain, g (measured in electrons per analogue-todigital unit, e - /ADU). The Gain is simply a multiplicative factor. Number of electrons = Pixel ADU value x Gain First astronomical use: TI 800 x 800 for HST (~ 1980)

16 How CCD s work (II) Reading out a CCD can take time, especially for the largest arrays with many pixels: the mode is different for video or astro The r.o.n. is ~ 1/ t You can use several amplifiers in parallel, but it is then difficult to have the same gain everywhere! In drift scanning CCDs are read out continuously (for example like the SDSS camera); no read-out time but integration time is fixed! Shuffling charge on CCDs can be a way to compensate for atmospheric seeing -- see the orthogonal charge transfer chips used in the Pann Starrs project

17 The positive points of CCD cameras... Extremely good quantum efficiency at redder wavebands close to 100%! Limiting magnitudes increased by four to five magnitudes! Sensitivity of telescopes is now limited by light collecting area and not the detector area. CCD s are linear over a very wide flux range: calibrations are easier! Output is digital!

18 Some problems with CCD detectors Poor blue (< ~4500 Å) response (too small penetration depth of blue photons) Megacam (Canada-France Hawaii Telescope) is one of the few wide-field detectors with good U* response Adequate surface treatment can help solve this problem, but at the cost of fringing! It s hard to make large monolithic CCD detectors: make mosaics from buttable detectors, but problem of curved focal planes. Charge transfer efficiency between wells is not 100% and degrades over time especially in space-based detectors. Sensitive to Cosmic Rays Cryogenic cooling (normally liquid nitrogen is used) is necessary to reduce dark current Reading out CCDs can take time!

19 Examples of unwanted effects... Diffraction spike (secondary mirror) Not specific to CCD s! Cosmic ray hits in megacam-u* image Charge transfer bleeding in STIS detector

20 Some real image defects (Hubble/WFPC2): Bleeding Ghost Cosmic ray

21 Charge Transfer (In)efficiency CCDs are read out by clocking charge along registers. These transfers are impeded by radiation damage to the chips. This effect gets worse with time and is worse in space, This image is from the STIS CCD on Hubble. Note the vertical tails on stars. Can degrade photometry and astrometry

22 More unwanted effects... Fringing : due to «thin» CCD... Stronger in red than in blue Dead columns Non uniform dark current (e.g. heating by the amplifier) All that can (needs to...!) be corrected during data reductions. (Flat Fielding!)

23 Geometric Distortion Cameras normally have some distortion, typically a few pixels towards the edges, It is important to understand and characterise it to allow it to be removed if necessary, particular when combining multiple images. Distortion may be a function of time, filter and colour. HST/ACS/WFC - a severe case of distortion - more than 200 pixels at the corners. Large skew. Due to optics, not CCD!

24 Characterising CCD detectors Quantum efficiency (QE): Can be as high as 90% in visible wavelengths for the best detectors today. Specific treatments to improve blue (or red) response. Resolution/pixel size: typical pixel sizes from 10 to 30 microns. For the images to be properly sampled, there must be at least two resolution elements per PSF (point spread function), and better three... otherwise the images are undersampled.

25 Characteristics of CCD s (II) Dynamic range/full well capacity/saturation level. Around 10 5 e -. Close to saturation level, non-linearity effects start to show up. It s important to choose an exposure time short enough so that objects of interest are not saturated! (But long enough so you are skynoise limited). Multiply exposures to be able to remove CR! Even in the absence of illumination, CCDs produce a thermal signal or dark current. Measure during bad weather...

26 Sources of noise in CCD observations Dark current Even in the absence of illumination, CCDs produce a signal or dark current. Typical values are 10 e - /hr if cooled to ~ -140 (LN) Readout noise This is an electronic noise (amplifier, or else), usually quoted in electrons per pixel. It does not depend on exposure time. Photon noise (object, or sky) Photon noise follows a Poissonian distribution (counting statistics): σ = n For bright objects, object photon noise dominates For faint objects, the photon noise from the sky dominates. Sky background increases at longer and longer wavelengths. In order to avoid being dominated by read noise, CCD exposures should be long enough to accumulate enough counts in the sky background to be sky noise dominated. Usually difficult in spectroscopy!!

27 Sky background Moonlight is yellow, so better to observe in redder bandpasses near the full moon Night sky from Paranal At longer wavelengths airglow lines cause significant variations in sky brightness

28 The CCD equation Object Signal: S obj = E * D 2 Δλ η τ t E in ergs/cm 2 /s/å or W/m 2 /Å, flux from object D Diameter of telescope, Δλ (bandpass), η (quantum efficiency), τ (transmission), t exposure time Total signal: S tot = (E * + (2)E s α 2 ) D 2 Δλητt + d ark. t + n 2 R 2 (α seeing, n number of pixels corresponding to α) Total counts per pixel, in electrons Astronomical Source Sky background Dark current Read noise N(oise) = square root of total signal Carefull! if calculation made in ADU or in electrons, or in photons... per pixel, or per total area of object!

29 A rough calculation of signal to noise Usually, Dark Current is negligible We assume the read-out noise is negligible also IF object brighter than sky, then: S/N = [ (E *. D 2 Δλητt ] 1/2 (~ independant of seeing) and Limiting magnitude E lim is ~ (S/N) 2 1/ [D 2 Δλητt] goes as 1/D 2 or 1/t for a given S/N IF object fainter than sky, AND r.o.n. negligible, then: S/N = E *. D/α. [Δλ η τ t /E c ] 1/2 or: Limiting magnitude: E * = S/N. α/d. [E S / (Δλ η τ t)] 1/2 Goes as 1/D, or 1/ t only and Seeing is as important as diameter of telescope...! This explains the performance of the HST despite its only 2.4m diameter.

30 How long should I integrate for?? You use the above equation...but you have to know the parameters!! Most observatories will offer exposure time calculators which will allow you to estimate how long to integrate to reach a given S/N But you need some critical eye...input parameters can be wrong, software have bugs, etc...so always have a rough check with the above equation yourself. But if you have (or can get) real data taken with the same instrumental setup, it s almost always better to use this, and to scale it up to your particular application. In practice, CCDs have finite full well capacity and the non-zero sky background means that integrating infinitely long is impractical (images would saturate on the sky). Also detectors have bad pixels / columns, so it is in general better to take several shorter exposures which are then dithered or offsetted on the sky and which can then be combined afterwards. This also helps to remove C.R. Important to choose large enough dithers to overcome detector defects. This is however only practicable for imaging, not always for spectroscopy...

31 CCDs and CCD mosaics

32 The next generation of large mosaic detectors...

33 { Electron Multiplying CCDs. Conventional CCD EMCCD Image Area Image Area (Architecture unchanged) On-Chip Amplifier On-Chip Amplifier On-Chip Amplifier Serial register Gain register Serial register The Gain Register can be added to any existing design

34 Normal Serial register Electron Multiplying CCDs. Image Area Store Area 1e - in Multiplication register Avalanche multiplication takes place in Multiplication Register, using an HV clock (40-45Volts). Multiplication register 520 multiplication stages 1000e - signal out Multiplication register Devices used : Standard MOSFET amplifier CCD x 128 CCD87/97 512x512 CCD201 1K x 1K

35 Multiplication Noise The EMCCD multiplication process increases the Poissonian noise in the image EMCCD Photon Transfer Plot

36 Electron Multiplying CCDs. Benefits : 1) Negligible read noise, therefore capable of seeing single photons 2) Read noise independent of read-out rate Disadvantages : 1) At present only available in small formats. 2) Suffer from Multiplication Noise that limits performance at high signals. Applications : Low signal regimes where detector read noise is the limiting factor. This could be an intrinsically faint source or a brighter source observed at very high frame rates. -> Adaptive optics wave front sensing -> High time resolution spectroscopy

37 QUCAM2 : an EMCCD camera on ISIS. 1k x 1k pixel image area. 3.2second full frame read out. Frame Transfer design

38 Superconductive Tunnel Junction ESA tests (M. Perryman, 2009) Charge proportionnal to the incoming photon energy R (spectral resolution) limited by the statistics of the generated electrons (R ~ E, a few 10) Working at very low temperature, in photon-counting mode Tests made at the WHT Developing a 2D array, and searching for higher T semiconductor material: S-Cam, 10 x 12, ESA 1m, R~15 R Photon Energy

39 Photonic Multi-object spectroscopy? Image courtesy of J. Allington-Smith (CfAI - U. of Durham) Use of photonic lanterns to transfer multi into single-- mode fibers (p=v²/4, V=kₒ a sinθ) (kₒ=2π/λ) 40

40 The photonic integrated multimode micro-spectrograph (PIMMS) The PIMMS instrument concept J. Bland-Hawthorn et al, Proc. SPIE 7735, 77350N (2010) 41

41 MKID s Optical/NIR Microwave Kinetic Inductance Detectors ARCONS 2024 pixel MKID array DARKNESS 10,000 pixel MKID array Rupert Dodkins University of Oxford

42 Operating Principle Incident photons briefly change the total inductance of superconductors through the Kinetic Inductance effect 1 MKID pixel (microresonator) Resonant frequency (and amplitude) shifts P. Day et al., Nature 2003 Optical/NIR MKIDs MKIDs background: 2/3

43 Specifics of IR detectors The thermal background could rapidly saturate the detector thus separate the detection area from the reading area. In the near-ir (λ <~ 2.5 µ), one uses Hg-Cd-Te (variable proportions) For longer wavelengths, material changes (InSb, SiGa, GeGa, etc ) In the thermal IR, necessity of rapid read-out, and wobbling (background subtraction: the background is much more important than the signal, and can change rapidly)

44 Hg (1-x) Cd (x)te: μ, 2k x 2k (e.g. VISTA) 20μ pix, 78 ᵒK μ, 18μ, 37 ᵒK, (JWST) InSb μ, 2k x 2k, 25μ, 32 ᵒK SiAs (IBC) 5-28μ, 1k x 1k, 18μ, 8 ᵒK, (WISE) Various IR detectors 5-28μ, 1k x 1k, 25μ, 7 ᵒK (JWST) MIRI 1k x 1k Hawaii 2k x 2k JWST Detectors

45 End of Part I Questions?

46 Part-II: Reducing CCD observations

47 CCD calibration frames Raw images from CCD detectors are not immediately usable for scientific exploitation but are instead contaminated by several instrumental effects. which need to be removed. Bias frame: There is a fixed offset level applied to all data. There may also be structure in the bias. It can be measured by examining the overscan region of each CCD, or take a zero-exposure (= bias) frame. Dark frame: Even without any signal, a dark current of a few electrons may be present. Ideally, one should take a dark frame of the same exposure time as the real data. Flat field: the pixel-to-pixel sensitivity of the detector is not constant. A uniformly (Flat...) illuminated exposure can measure this. Flat-field is wavelength dependent. Fringe frame: Thinned CCDs do produce interference fringes at longer wavelengths. The amplitude and separation of the fringes can be time-dependent, making it difficult to remove them (night-sky variations...) Megacam data with badly removed fringe pattern

48 Subtract bias or overscan Typical CCD pre-reduction steps Master bias can be computed from median of many frames. The overscan region can be used to compute the bias level Subtract dark frame Master dark frame should be computed from many observations of several tens of minutes scaled up to the appropriate exposure time Divide by twilight or dome flat-field (filter dependent; imaging only) Dome flat-fields are almost always not flat enough for wide-field cameras. Sky flats are time limited...the best is to use super-flats (imaging only) Make sure that there are enough counts in the flat-fields so they are effectively sky-noise limited. Subtract sky flat / fringe frame At each stage in the reductions you should verify that the noise per pixel in the individual frames decreases. It should approach the Poisson limit. If the reduction steps add noise, then you are doing something wrong!

49 Other calibration frames you may need (for imaging) A bad pixel mask is computed from a single image and identifies where the stuck pixels, bad columns and cosmic rays and satellite trails are in each image. A weight map or coverage map which is computed from the normalised flat-field. The weight map is necessary to compute magnitude errors and detection thresholds correctly in a stacked image. Weight maps are less important for single detector cameras. Not all sections of the image will receive equal exposure time (vignetting..., shutter problems, etc...) Megacam single image weight map (flat-field + bad pixel mask + cosmic ray mask)

50 FITS: the standard for astronomical images FITS=flexible image transport system FITS is an architecture independent way to store and transfer astronomical data Simplest FITS file consists of a primary header comprising of n(36x80) ASCII characters followed by binary data FITS files can also contain tables or spectra. In multi-extension-fits there are several extensions, which can be useful for mosaic camera data. topcat, ds9, fv can be used to manipulate and inspect fits tables

51 Complete reduction step for CCD imaging data Carry out the pre-reductions described previously. If you have a series of images on the same part of the sky you will need to separately flux calibrate and astrometrically calibrate them. Images which are dithered around the sky will need to be re-aligned. For small detectors, linear shifts are sufficient; for wide-field imaging detectors interpolation and reprojection will have to be carried out Next the images need to be combined to produce a single, stacked image. The last step involves the extraction of catalogues. Spectroscopy: next talk!! CCDPROC (iraf) SCAMP (Terapix) SWARP (Terapix) or IMSHIFT (IRAF) SWARP (Terapix) or IMCOMBINE (IRAF) Sextractor or PHOT