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1 Astronomical Observational Techniques and Instrumentation RIT Course Number 1060-771 Professor Don Figer Spectral resolution, wavelength coverage, the.

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Presentation on theme: "1 Astronomical Observational Techniques and Instrumentation RIT Course Number 1060-771 Professor Don Figer Spectral resolution, wavelength coverage, the."— Presentation transcript:

1 1 Astronomical Observational Techniques and Instrumentation RIT Course Number 1060-771 Professor Don Figer Spectral resolution, wavelength coverage, the atmosphere and background sources

2 2 Aims and outline for this lecture describe system requirements for spectral resolution and wavelength coverage describe atmospheric effects on observations summarize primary background sources and effects on observations

3 3 Need for Spectral Information “A picture is worth a thousand words.” --Barnard “A spectrum is worth a thousand pictures.” --an astronomer A spectrum is the distribution of flux versus wavelength. Spectra are critical for making measurements of physical properties of astronomical objects. –low resolution spectra reveals spectral energy distribution (SED) –high resolution spectra reveals emission/absorption features

4 4 Spectrum of Arcturus

5 5 High Resolution Spectroscopy High resolution spectroscopy can yield –velocity –abundances –temperature –virial mass (through integrated light) –molecular excitation –interstellar absorption –ionization state –others The example on the right shows high resolution infrared spectra of red supergiants. Note the many CO rotational-vibrational absorption lines. The slight translations in the wavelength axis from star to star shows velocity differences. Weaker features to the left reflect atomic absorption and can be used for abundance analysis. velocity difference CO bandhead atomic absorption

6 6 Low Resolution Spectroscopy Low resolution spectroscopy can yield –emission mechanism (via spectral index) –temperature –molecular content –interstellar molecular absorption The example on the right shows low resolution infrared spectra of red supergiants. Note the CO bandhead. The depth of the bandhead indicates that these stars are cool. For this application, it might be just as well to obtain a high resolution spectrum over this waveband. In some cases, a low resolution spectrum can reveal curvature over a broad range of wavelengths – something difficult to detect in high resolution spectra. CO bandhead

7 7 Spectrograph/spectrometer A spectrograph/spectrometer is an instrument that can measure intensity versus wavelength. –dispersive prism grating –nondispersive filters, circular variable filter (CVF) Fabry-Perot

8 8 Spectroscopy System Design Criteria Efficiency: amount of light directed into the spectrum Resolution: minimum spatial and spectral separation between resolved features Scattering: light directed out of the spectrum or in undesired locations in the spectrum Stability and calibration accuracy: ability of system to maintain or reproduce wavelength versus pixel relationship Background: light from all sources other than target, might be sky, telescope, optics, baffles, etc. Packaging: compactness of layout

9 9 Spectrograph Cartoon

10 10 Continuum and Emission Lines

11 11 Absorption Lines

12 Equivalent Width Equivalent width is a measure of flux in a line. It is independent of instrument parameters. 12

13 13 Spectroscopy System Design Form Example graphic shows simple layout for spectrograph grating normal grating collimator slit telescope detector lens

14 14 Spectral Resolution Spectral resolution is a measure of the ability to separate nearby features in wavelength space. Delta lambda –often set to the full-width at half-maximum of an unresolved line –can be measured in the data –depends on data analysis –can be limited by diffraction, slit width, detector sampling

15 15 Spectral Resolution: Prism A prism disperses light with a transmissive optic that has chromatic index of refraction.

16 16 Gratings: Grating Equation The grating equation gives the geometry for constructive interference between facets of a grating. The equation can be derived by setting the optical path difference between light scattering from two adjacent facets to an integer multiple of the wavelength. In more standard nomenclature, where, m is the order number and T is the groove density. Note that the equation reduces to the law of reflection for m=0.

17 17 Gratings: Angular Resolution Angular resolution describes the variation in output angle versus wavelength. It can be found by differentiating the grating equation. Note that the angular resolution can be increased by increasing the output angle. For a fixed wavelength, angular resolution is dependent on geometry.

18 18 Gratings: Diffraction-limited Resolution The diffraction-limited resolution of a grating is simply Nm, where N is the number of illuminated grooves and m is the order number.

19 19 Gratings: Slit-limited Resolution The resolution, or resolving power, is usually limited by slit width. (“primed” quantity is at image plane)

20 20 Gratings: Diffraction-limited Image Imagine that we set the slit width to be equal to the full width at half maximum of a diffraction-limited spot. (Assume Littrow). Note that the resolution can be increased by increasing the size of the collimated beam (and size of grating), increasing the output angle, or decreasing the wavelength (assuming that the decrease in wavelength is accommodated by a decrease in the spot size and decrease in slit width).

21 21 Gratings: Increasing Resolution From the previous equations, it appears that resolution can be increased by: –decreasing the slit width (which might decrease the transmitted flux) –increase the collimator focal length (which will increase the beam size and required size of the grating) –increase the output angle (which will increase the required size of the grating) There is another possibility that is not apparent from the grating equation, given that its derivation assumed that the index of refraction of the material surrounding the grating was one. A more general derivation would show that the realized angular dispersion is multiplied by n. –This effect is used in an immersion grating –Silicon is a particularly useful material for this purpose because it has a high index (n~4), although it is only transmissive in the infrared.

22 22 Gratings: Immersion Grating Resolution The resolution can be increased by immersing the grating surface in a high index material. Note that angular dispersion is n times greater than for case without immersion.

23 23 Gratings: Echelle Grating Angular dispersion is dramatically increased when using an “echelle” grating. In this case, the output angle is typically >60 degrees. Because an echelle is used at a large angle, it must be rectangular. The grating facets must be ruled at a large angle, because: –In order to maintain high efficiency, the input and output angles should be the same (“Littrow condition”). –In order to provide high resolution, the output angle must be large.

24 24 Gratings: Echellogram HIRES order format showing the solar spectrum (note that color was artificially added for illustrative purposes) + +

25 25 Gratings: Immersion Grating

26 26 Fabry-Perot Etalon A Fabry-Perot etalon is a filter that selectively transmits a narrow range of wavelengths. Like a thin film, it transmits the most flux at wavelengths that satisfy the thin film condition. Finesse is a function of reflectivity.

27 27 Wavelength Coverage: Free Spectral Range Multiple orders of the same wavelength range are dispersed to either side of the order of interest. The non-overlapping wavelength region in the order of interest is defined as the “free spectral range.” Outside of this range, light of different wavelengths in multiple orders overlap. mm+1m-1 order=m 1 2 1 2 FSR overlap FSR= /m

28 28 Atmospheric effects Absorption –reduced flux from source –difficult calibrations Emission –increased background noise –reduced integration times –difficult calibrations (subtracting time-varying components) Turbulence –increased object size (“seeing”) All effects vary with wavelength, time, altitude, line-of-sight

29 29 Atmospheric absorption Molecules are the dominant absorbers (H 2 O, CO 2, O x ) Strong function of: –wavelength, atmospheric wavebands –time, frequent calibration (hour timescales) –altitude, mountain-top observing sites –line-of-sight, limited target access and frequent calibration

30 30 Atmospheric absorption versus Sharp cutoffs –defined primarily by H 2 O –shape wavebands Higher transmission between lines with higher resolution Can introduce large calibration errors for low resolution observations (MNRAS, 1994, 266, 497)

31 31 Wavelength Regimes: near/mid-infrared: Atmospheric Transmission

32 32 Wavelength Regimes: near-infrared: Atmospheric Transmission JHKLM

33 33 Atmospheric absorption versus  - high res Array defects CO 2 absorption lines Keck II 10-m  + + R =  

34 34 Particle number densities (n) for most absorbers fall off rapidly with increasing altitude. Density of atmospheric constituents modeled as exponential. x 0,H 2 0  km, x 0,CO 2  km, x 0,O 3  km So, 95% of atmospheric water vapor is below the altitude of Mauna Kea. Atmospheric absorption versus altitude

35 35 Atmospheric absorption versus airmass The amount of absorbed radiation depends upon the number of absorbers along the line of sight Atmosphere AM=1 AM=2

36 36 Atmospheric emission Blackbody (thermal) Molecular (OH)

37 37 Thermal emission  spectral radiance, brightness, specific intensity: I =  cos  B (T) W m -2 Hz -1 sr -1 B (T) = 2h 3 1 c 2 exp(h /kT) - 1 Planck (blackbody) function:   emissivity (dimensionless) max (  m) = 3674 K T Peak in B or B : B (T) = 2hc 2 1 5 exp(hc/ kT) - 1 max ~10  m for room temp.

38 38 Total power onto a detector: P =  A  sky  B (T sky )  : transmission of all optics x Q.E.  sky : emissivity of sky A: telescope area  : solid angle subtended by focal plane aperture  : bandwidth B (T sky ): Planck function At 10  m, typically:  ~ 0.2,  ~ 0.1, A  ~ 3x10 -10 m 2 Sr  ~ 1.5 x 10 13 Hz (10  m filter), T ~ 270 K P ~ 10 -9 W or ~ 4 x 10 10  s -1 Atmospheric emission: Blackbody

39 39 Atmospheric emission: A Note on Solid Angle Background from the sky can be calculated by multiplying the blackbody equation by the area of the sky seen by the detector and the solid angle subtended by the telescope aperture. The A  product is invariant in a system. A telescope A sky detector D x d F   It is easier to use area of telescope and solid angle of sky.

40 40 Atmospheric emission: OH lines OH lines Photons s -1 m -2 asec -2  m -1 Wavelength {  m}

41 41 Array defects OH lines + + R =   Atmospheric emission: Molecular lines Array defects Keck II 10-m + + R =  

42 42 Atmospheric emission versus time

43 43 Atmospheric Turbulence Static atmosphere bends light (n=1.000273 at 2  m) Dynamic atmosphere distorts the wavefront Data from the William Herschel 4.2 metre Telescope in La Palma. Wavelength 689.3nm, Exposure time 30ms per frame. The data were taken with a simple Speckle Camera situated on the GHRIL platform of the telescope. The data sets were taken with the Speckle Imaging group at Cardiff University.  Orionis (“Betelguese”)

44 44 Atmospheric Turbulence A diffraction-limited point spread function (PSF) has a full-width at half-maximum (FWHM) of: In reality, atmospheric turbulence smears the image: At Mauna Kea, r 0 =0.2 m at 0.5  m. “Isoplanatic patch” is area on sky over which phase is relatively constant.

45 45 Lick 3-m (1994)

46 46 Keck 10-m (1997)

47 47 HST/NICMOS (1997)

48 48 VLT/AO (2002)

49 49 Keck/LGSAO (2006)

50 50 Atmosphere –thermal –molecular Telescope –thermal –scattering Zodiacal light Astronomical sources Background - sources

51 51 Thermal OH –The average OH line intensity is approximately 25,000  s -1 m -2 asec -2  m -1. See Maihara et al. 1993. –The continuum between lines is about 50 times lower than this value (in the H band). Background - sources: Atmosphere

52 52 Mirrors Baffle edges and walls Secondary support Background - sources: Telescope - scattering

53 53 Background - sources: Astronomical Astronomical objects can be objects of interest or noise contributors, depending on the project. –Sunlight, moonlight –Light scattered by solar system dust (“zodiacal”) –Light emitted (thermal) by solar system dust (“zodiacal”) –Stars (especially in a crowded field) –Light emitted by interstellar dust (“cirrus”)

54 54 Background - sources: Astronomical

55 55 Background - sources: COBE data

56 56 Background noise is the quadrature sum of shot noise for each background source. Background and Signal-to-Noise Ratio Recall the expression for SNR. The noise is the quadrature sum of uncorrelated noise sources.

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