Instrumentation Concepts Ground-based Optical Telescopes

Aug-Nov, 2008 IAG/USP (Keith Taylor)

Instrumentation Concepts

Ground-based Optical Telescopes

Keith Taylor(IAG/USP)

Aug-Nov, 2008

Aug-Sep, 2008 IAG-USP (Keith Taylor)


CCDs in AstronomyCCDs in Astronomy& Signal to Noise

Ratios

Charged coupled devices (CCDs) have been used in astronomy since the late 1970s. They are now nearly ubiquitous in observations made at wavelengths between the near-IR(~1µm) and the X-ray. They have transformed the way astronomy is done.


Basic CCD Design Apart from sensitivity, the key design issues for solid

state arrays are to localize photon-produced charges on their surfaces and then arrange to amplify and read them out without distorting the image or introducing unacceptable amounts of noise.

A CCD is a charge-transfer device. Its storage, transfer, and amplification electronics are all fabricated on a single piece of silicon (unlike most IR arrays). During an exposure, it traps released photoelectrons in small voltage wells. After the exposure, the electrons are shifted in a series of "charge-coupled" steps across the CCD surface, amplified, read out of the CCD, and stored in a computer memory. This is "destructive readout"---i.e. one cannot read the same

signal again (unlike other array architectures where each pixel is coupled to a separate amplifier).


Basic Structural Element

The basic element in a CCD design is a "Metal-Oxide-Semiconductor" capacitor. This serves both to store photoelectrons and to shift them wholesale. The bulk material is p-silicon on which an insulating layer of silicon-oxide has been grown. P-silicon can be made to have very small numbers of free electrons ("high resistivity") before exposure to light; this is important for best performance. A set of thin conducting electrodes of semitransparent "polysilicon" are applied.

Before exposure, the central electrode is set to a positive bias while the two flanking electrodes are set negative. This creates a "depletion" region under the central electrode containing no holes but a deep potential well to trap electrons. The region shown is about 10µm thick.


Operation Sequence During exposure (controlled by a

separate shutter), light enters through the "front-side" electrodes. Photoelectrons generated under the central electrode will be attracted toward the electrode and held below it. The corresponding holes will be swept away into the bulk silicon. Good performance requires little diffusion of electrons away from the potential well.

The surface of the CCD is covered with MOS capactitors in a pattern like that at left. In this particular design, there are three electrodes per pixel. A single pixel is shown shaded in the diagram. Typical pixel sizes are 10-40µm. The "parallel shift" registers are shown as rows running across the whole face of the CCD. These are separated by insulating "channel stops."


Clocking At one end is a column of "serial

shift" electronics and an output amplifier. Note that there is only one amplifier in this design. Contemporary large chip designs involve several amplifiers (but always many fewer than the number of pixels!).

At the end of the exposure, readout of the collected electrons is accomplished by cycling ("clocking") the voltages on the electrodes such that the charge is shifted bodily to the right along the rows. The figure on the left shows how this is done. Good performance depends on near-100% transfer of the electrons to/from each electrode with no smearing or generation of spurious electrons.

Each parallel transfer places the contents of one pixel in each row into the serial register column. This column is then shifted out vertically through the output amplifier and into computer memory before the next parallel transfer occurs.


Serial and Parallel Registers


ADU Conversion For storage in memory, the electrical signal generated by the

amplifier must be digitized. This is done by an "analog-to-digital converter“ (or ADC). This is normally adjusted such that one digital unit corresponds to more than one photo-electron. Typical values of this conversion are 2 to 8 electrons per stored digital unit.

The stored values are called "ADUs", for analog-to-digital-units. The corresponding constant of transformation, normally quoted in units of "electrons per ADU", is often called the "Gain" (although this is confusing nomenclature because a larger Gain results in reduced ADU values).

Note that the use of such a conversion importantly affects the statistical properties of the recorded signal. If x is the recorded signal in ADU's, y is the original signal in photo-electrons, and G is the gain, then we see that:

2(x) = 2(y)/G2


CCD Design Issues

CCDs have undergone a long optimization process since 1980. Contemporary designs have excellent performance but still require careful calibration in order to overcome inherent limitations. There is also only a handful of reliable manufacturers of professional-grade chips.

Here are some of the issues affecting electronic design and manufacture of CCDs:


Poor Blue Response (< 450nm)

Mitigation: Use special thin, polysilicon material

for electrodes. But cannot be completely transparent.

Special Coatings: "Anti-reflection" coatings trap photons, causing multiple reflections as in Fabry-Perot etalons, and therefore enhance absorption. "Downconverter" coatings are phosphors which absorb blue photons but emit green photons at wavelengths where the CCD QE is higher (e.g. "mouse milk," coronene, lumogen).

"Thinned, Backside-Illuminated" chips: shave off silicon subtrate, leaving only ~10µm deep unit; illuminate from backside; greatly improves blue response.

Caused by absorption in bulk Si and by electrode structures in "Frontside-Illuminated" chips.


Difficulties with Thinned Chips

Fragile

Non-uniform thinning

Surface trapping by SiO2 layer of photo-electrons produced nearby (shorter wavelengths)

Interference effects if wavelength ~ chip thickness (i.e. in IR). Strong spatial modulation of response = "Fringing". Especially serious for night sky emission lines. Can be well calibrated for narrow-band filters or for broad-band filters. Hard for intermediate band filters.

Thinning reduces red response. For good response 500—1000nm, earlier generations preferred thick

(~500µm) front-illuminated chips. Now use of deep (or fully) depleted chips is preferred.


CCD Fringing


Charge Transfer Efficiency (CTE)

CTE can be better than 99.999% per transfer, but has to be, since throughput of chip with 2048 required shifts = CTE2048 ~0.98

Radiation damage to CCD's in space seriously decreases CTE over several years' time.

Mitigation: Operational: add (electronically or with diffuse light source) a

pedestal background signal (a "fat zero") over whole chip to increase mean electron density per pixel. However, adds additional noise and not suitable for very faint source applications.

Technical design: change number phases, clocking cycles; add "minichannels."


Advantages of CCD Detectors(Quantum Efficiency)


CCD High Quantum Efficiency

This had tremendous impact on astronomical imaging & spectroscopy. It meant the detection threshold with any instrument was extended 4-5 magnitudes over film and that therefore a 1-m telescope could now pursue the kind of science previously possible only with 4-m class telescopes.

A key corollary: since we are already near 100% QE, at least in the optical region, achieving significantly lower threshold fluxes requires larger telescopes rather than better detectors.

NB: "Quantum yield" can be over 100% for far-UV and X-ray photons (i.e. more than one photoelectron can be generated but fewer than 100% of photons produce responses).

To 80-90% at peak in optical. Much effort was expended to reach these high levels.


CCD Characteristics LINEAR RESPONSE: Until approach full-well capacity (typically 200,000 e/pixel).

This implied much improved flux calibrations and much higher precision for flux measurements at all levels.

EXCELLENT DYNAMIC RANGE: Typically 104.

WIDE WAVELENGTH RESPONSE : Intrinsically sensitive from X-ray to ~1µm. Other materials (e.g. InSb2, HgCdTe) with similar architectures are usable in the IR.

GEOMETRICALLY STABLE: good for astrometry

ONLY LOW VOLTAGES REQUIRED (~5-15 v)

RELATIVELY CHEAP, AVAILABLE, SIMPLE: Compared to other digital 2-D systems. Standard chips cost ~$200K.

RELATIVELY LOW NOISE : Compared to many other classes of astronomical detectors, e.g. Pg plates, Reticons, SEC vidicons, etc. But noise is not negligible. Typical read noise now 2-5 e/pixel, and dark current is largely suppressed by cooling.

SMALL PIXELS : Typically 10-30µm. Usually an advantage, but should match pixel size to 1/2 of smallest resolvable picture element in optical system.


Special Format/Readout Designs

Rapid clocking/video: inherent in CCDs intended for TV application. For bright sources, readout rates of 100 MHz are possible.

Drift Scanning: Basic technique is to clock the chip slowly along columns while

moving the telescope to keep a target centered on the moving electron cloud. Continuously read out the chip to produce a strip-image of the sky. The integration time is set by the drift time across the chip.

The simplest approach is to align the CCD columns east-west, keep the telescope fixed on the sky, and clock the chip westward at the sidereal rate.

Drift-scanning is now a standard method for wide-field CCD sky surveys, most notably the Sloan Digital Sky Survey

By changing electrode structure and clocking cycles, one can arrange for many different integrate/readout modes.


Continuous Drift-scan


Drift-scanned Image of a Galaxy

Messier 51, The Whirlpool Galaxy. The SDSS image of this famous spiral galaxy (interacting with a smaller neighbor at the lower left) occupies about three one-millionths of the total sky area imaged by the SDSS.


Spectroscopic Survey

Slices through the SDSS 3-dimensional map of the distribution of galaxies. Earth is at the center, and each point represents a galaxy, typically containing about 100 billion stars. Galaxies are colored according to the ages of their stars, with the redder, more strongly clustered points showing galaxies that are made of older stars.

The region between the wedges was not mapped by the SDSS because dust in our own Milky Way.


A map of stars in the outer regions of the Milky Way Galaxy, derived from the SDSS images of the northern sky, shown in a Mercator-like projection. The color indicates the distance of the stars, while the intensity indicates the density of stars on the sky. Structures visible in this map include streams of stars torn from the Sagittarius dwarf galaxy, a smaller 'orphan' stream crossing the Sagittarius streams, the 'Monoceros Ring' that encircles the Milky Way disk, trails of stars being stripped from the globular cluster Palomar 5, and excesses of stars found towards the constellations Virgo and Hercules. Circles enclose new Milky Way companions discovered by the SDSS; two of these are faint globular star clusters, while the others are faint dwarf galaxies.


A mosaic showing 36 of the the 500+ Type Ia supernovae discovered by the Sloan Supernova Survey. Each image is centered on the supernova, which usually stands out as a bright point near or within the galaxy that hosts it. The light of the supernova, powered by the thermonuclear explosion of a single white dwarf star, can outshine that of the tens of billions of stars in its host galaxy. Type Ia supernovae have a constant intrinsic luminosity (after a correction based on the time over which their light rises and falls), so their apparent brightness can be used to infer their distance. The primary goal of the Sloan Supernova Survey was to measure the expansion of the universe with high precision over the last four billion years of cosmic history, to help understand why that expansion is speeding up over time despite the decelerating gravitational effect of atoms and dark matter.


In addition to images, the SDSS has measured the spectra of light from more than a million celestial sources.

This picture shows the spectra of more than 46,000 quasars from the SDSS 3rd data release; each spectrum has been converted to a single horizontal line, and they are stacked one above the other with the closest quasars at the bottom and the most distant quasars at the top. Bright bands show the emission produced by specific ions of hydrogen, carbon, oxygen, magnesium, and iron. For more distant quasars, these emission lines are shifted to longer wavelengths by the expansion of the universe. This redshift of spectral lines is what is used to determine the distances to quasars and galaxies.


The 2.5-meter SDSS telescope at Apache Point Observatory.

All SDSS imaging and spectroscopy were carried out with the Sloan Telescope, equipped with a large format digital camera and fiber-fed spectrographs that measured spectra of 640 objects at a time.


Nod and Shuffle (N&S)(“Va et Vent”)

“Nod and shuffle” (N&S) is a mode of observing where parts of the charge-coupled device (CCD) are used as a “storage register” in a beam-switched image. Beam-switching is achieved by rapid alternation between object and sky positions (“nodding”), which is undertaken with no readout penalty. Instead, the sky image is shuffled to a storage region.

Typically, nodding takes place every 30–60 seconds, which is a timescale faster than the variations of airglow emission lines. Because both the sky and objects are observed quasi-simultaneously through the same optical path, slits and pixels, N&S provides an order of magnitude improvement in sky subtraction opening up significant new observational capabilities for large telescopes. For example, very deep integrations (ten times longer than ispractical with conventional spectroscopy) are possible with N&S at very high slit densities.

"Nod and Shuffle": this technique takes advantage of the capability to shift an image wholesale on the CCD without reading out to obtain much better sky subtraction (e.g. in near-IR where atmospheric OH emission is very bright and variable).


N&S – how?

A

B

Night sky emission lines

Nod distance on sky

Shuffle distance on CCD

A-B

Schematic of nod-and-shuffle for a single slit. When the telescope is in the “object” position, CCD area “A” records a spectrum. The “sky” position records the nodded spectrum (in this case the telescope has been nodded a few arcseconds along the slit direction).

The area “B” is non-illuminated by the mask and serves as a storage area for the “sky position.” The image difference subtracts the sky, and leaves a positive and negative object spectrum for subsequentextraction.

Target


Illustration of Masks with Different Shuffle

Distances.Top row shows the input masks; the bottom row shows the resulting shuffled object-sky spectra. The black spectra correspond to image A, and the light grey spectra correspond to image B, as before;

(a) a mask where the shuffle is only a few pixels, and the sky is stored below the object – appropriate for extended, relatively low-source-density regions

(b) a mask where the shuffle is large, appropriate for cases with compact regions of high-source density

(c) an intermediate case. This compromise has the advantage of allowing a high-density extended field to be tacked while minimizing the number of object-storage interfaces where it is necessary to leave gaps

Note that in reality the area we have shown as a single “detector” is three CCDs.

Shu

ffle

d S

pect

raS

lit M

asks

Buffer storage Buffer storage Buffer storage

(a) (b) (c)


On chip binning Change clocking to combine electrons from several

pixels together before reading out through amplifier. E.g. combine a 2x2 pixel region on the chip into a single output pixel. This reduces the effect of amplifier readout noise on each pixel in the final data. Also reduces memory and storage requirements. Adds

considerable practical flexibility to CCD systems. Obviously, however, reduces the resolution of the output image.

Is useful for applications such as imagery of very faint, extended sources (e.g. galaxy halos), low spatial resolution spectroscopy, photometry of point sources under poor seeing conditions, etc.


On-chip binning(Illustrated)


Orthogonal Transfer CCDs (OTA)

Pan-STARRS project: Each focal plane consists of an 8 x 8 array OTAs. An OTA is a single silicon chip which itself consists of an 8 x 8 array of individually addressable 512 x 512 Othogonal Transfer CCD devices. Hence the total number of pixels per camera is (8x8)x (8x8)x(512x512) - about one billion.

Tip-tilt correction

over a very large FoV.

64 independent cells of TT


Other advantages of CCDs

UBIQUITOUS: Now almost universally used in astronomy (amateur & pro). Photographic materials and older electronic detectors being phased out.

IMMEDIATE DIGITAL CONVERSION OF DATA: The other advantages of array detectors notwithstanding, it is the immediate conversion of astronomical signals into a form capable of computer storage/analysis which has so dramatically transformed UVOIR astronomy since 1980. The practice, pace, & scope of UVOIR astronomy are entirely different now than in the "photographic" era that preceded widespread use of CCDs and other array devices. Digital conversion of images and spectra has enabled quantitative analysis of observations on a scale not possible before.

Among other things, rapid digital processing allows much improved use of observing time---notably in surveys (e.g. for variable sources in MACHO and supernova searches; moving targets such as asteroids/Kuiper Belt objects; or combined imaging/spectroscopic surveys such as SDSS, 2dF).


Limitations/Disadvantages of CCDs

SMALL SIZE : Individual chips typically less than 7 cm square. Cover only ~20% of high quality imaging field on typical modern telescope. Size limited by small fabrication yield of large, defect-free chips. Largest routinely available chips are 4096x2048.

However, MOSAICKING technology now well developed. Typical mosaic now uses 4096x2048 chips.


Examples of Mosaics Sloan Digital Sky Survey: 54 chip mosaic for drift scanning in 5 bands simultaneously


Canada-France-Hawaii-Telescope MegaPrime mosaic: 40 4096x2048 chips covering 1 degree FoV


More Limitations/Disadvantages

CRYOGENIC COOLING REQUIRED: To reduce dark noise, cooling to below -100 C necessary. Thermoelectric coolers usually not adequate, so cryogens (e.g. liquid N2) required. Introduces many practical complications. Optimal low-T chips differ from commercial types used at room temperature in digital cameras, etc.

READOUT NOISE: Produced by variations in amplifier gain. Much effort to reduce. Now typically 2-5 e/pix (rms).

What matters is not the noise per pixel but rather the total noise per image area, which can extend over many pixels, depending on the application.

Even at these very low levels, RoN can compromise some types of observations (spectroscopy, surface photometry)

Can reduce RON effects in some applications using on-chip binning. In some cases, it may be more advantageous to use "pulse counting" detectors, which can unambiguously detect individual photons, than CCDs.

FULL WELL CONSTRAINTS: Bright sources which over-fill pixels can produce "blooming" or "bleeding" along columns, making parts of the chip other than the immediate vicinity of the source useless.



Response nonuniformities ("flat field" effects): Caused by small variations in masks used to manufacture chips, deposition irregularities, thinning variations, etc. Typically 5% pixel-to-pixel. Color dependent. Requires extensive calibration, with color-matching to targets. Use "dome flats" or "twilight flats." Special observing procedures to reduce flat field effects include:

Drift Scanning: as above Multiple Offset: Break exposure into 4-5 parts, offset ~50 pixels between

exposures. Combine exposures during data reduction. "Dithering" is a related technique involving smaller offsets to achieve sub-pixel spatial resolution. These methods result in reduced field of view because not all parts of the original field will have uniform exposures.

Sensitivity to cosmic rays: Especially thick chips. Must remove effects during processing. CR's are a major problem for spacecraft detectors (e.g. HST/WFPC2). Requires that exposures be broken into multiple parts (called "CRSPLITs" on HST) so that CR events can be detected. CR hits can be removed from data frames, but this always leaves "holes" which have less exposure than other parts of the image.

Sensitivity to x-rays: An advantage for X-ray astronomy, but some materials in vicinity of CCD's, e.g. special glasses used for windows, can produce a diffuse background of X-rays which add noise to observations.



Charge transfer efficiency: Good CTE is often possible only for signal levels above threshold ~10-50 e/pix. For low light levels, require adding "fat zero" (typ. 1000 e/pix) electronically or by preflash. Creates added noise (10-30 e/pix).

Availability hostage to commercial market: CCD availability has always been driven more by commercial and military applications than science. Scientific CCD manufacture represents only about 10% of the overall $1B CCD market. A serious near-term issue, since industry is moving to "active pixel sensors," a different technology with amplifiers incorporated in each pixel. Requirements for good performance at low (astronomical) light levels are considerably different than for room-temperature, short-exposure, mass-produced equipment.

Data glut! CCDs typically produce 2 x Npix bytes of data per readout, where Npix is the number of pixels. For a 8096x8096 mosaic, this is 130 MB. Data storage/manipulation was a serious problem when CCDs were first introduced, and this influenced the style of data processing systems such as IRAF. Disk and tape storage are now much improved, but long-term stability of massive amounts of data now being produced is a non-trivial issue. All that data is really important and worth saving. Or should we just leave it in th sky where it’s safe?


CCD Calibration Procedure

Data required Bias frames Dark frames (optional) Flat field frames Sky flat/fringe frames (optional) Flux calibrator fields Target frames


Subtract Bias frames Bias frames (zero exposure time) taken with chip

unexposed to light from telescope. These measure electronic pedestal of chip. For high precision, average many bias frames, before and after observations. Check for bias drift during night. With some chips, can determine electronic bias level from overscan region on chip. For optical "fat zero" subtract average of many frames.

Optional: Subtract Dark Frame: Median filter many long (~30 min) dark exposures; note possible LED activity of CCD electrodes. Scale result to integration time of each data frame before subtracting.


Divide by Twilight or Dome Flat-Field

To remove residual pixel-to-pixel variations (typical 5%). Make exposures of twilight sky (good diffuse source, but tricky to get exposure level right; only 2 chances per night). Or make many exposures of diffusely illuminated screen in dome (disadvantage: often not uniformly illuminated). Must be done for each filter used (color sensitive effects).

Note: photon statistics must yield S/N significantly in excess of final desired S/N---i.e. if desire S/N = 100, require over 10,000 photons per pixel in net flat field image.

Optional: Subtract sky flat/fringe frame: Remove night sky emission line fringing effects (worse in near-IR) by observing uncrowded field in night sky. Take several exposures, moving telescope by, say, 25 arc-sec between them. Use "adaptive modal filter" technique to zap star images and create mean sky frame. Scale to target frames and subtract. Resulting flat field as good as 0.1% can result. Require a sky flat determined this way for each filter.


Further calibrations Use multiple offsets, dithering or drift scanning for target frames: For

faintest possible photometry, use multiple-offset exposure technique---e.g. 500 sec exposures shifted by about 20 arc-sec each--- to reduce flat field errors. Always want more than one exposure anyway for "reality check," empirical determination of photometric errors, cosmic ray removal, etc.

Interpolate over or mask out bad pixels/remove cosmic ray events

Register and combine target frames: Re-register offset frames to sub-pixel accuracy (e.g. "Drizzle" software). Median filter. Trim off "lost" parts of image.

Calibrate flux scale: Observe "standard stars" in recommended CCD calibrator fields (e.g. star clusters) to determine nightly atmospheric extinction and telescope throughput. No need to take on non-photometric nights (variable clouds), but no flux calibration either. Important that calibrator stars overlap targets of interest in color. These frames are reduced exactly like target frames, apparent fluxes extracted, and correction factors to determine true fluxes of targets are obtained.


SCIENTIFIC APPLICATIONS

Sub-pixel registration of dithered HST imaging using "drizzling" technique. On left is one original frame (I-band, 2400 sec exposure). On right is result of drizzled processing of 12 such frames. Combined image has limiting magnitude of I ~ 28. It would be impossible to reach such levels with photographic detectors.


Example of improvement in color-magnitude diagram for star cluster 47 Tuc. On left is 1977 state-of-the-art CMD based on widefield photographic images with photoelectric calibrations (Hesser & Hartwick, 1977). On right is a 1987 CMD derived from CCD photometry (Hesser et al. 1987). The greatly improved photometric precision reveals new features of astrophysical interest: e.g. the thinness of the main sequence near turnoff, which places strong limits on the range of ages present in the cluser

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Instrumentation Concepts Ground-based Optical Telescopes