Astronomical Spectroscopy. Michael Cushing

Transcription

1 Astronomical Spectroscopy Michael Cushing REU Presentation June, 08, 2009

2 What Is a Spectrum? A stars have Teff ~10 4 K. Continuum H Line Absorption Jacoby et al. (1984, ApJS, 56, 257)

3 What is a Spectrum? SMA (865 µm) of Orion KL. Line Emission Beuther et al. (2005, ApJ, 632, 355)

4 Cartoon Diagram of a Spectrum f ( Intensity) Δλ resolution Wavelength or Frequency

5 Resolving Power Resolution: The difference (Δλ) between two wavelengths that can just barely be resolved. Rayleigh s criterion Resolving Power: R / / = v/c v = c/r

6 Resolving Power R ~ , SEDs, solid state and molecular bands R ~ 1000 (300 km s-1 ) Atomic and molecular features, red shifts R ~ 104 (30 km s -1 ) line profiles, radial velocities of stars, stellar abundances R ~ 105 (3 km s -1 ) ISM lines, molecular spectroscopy

7 R ~ An R=150 spectrum of a brown dwarf obtained with SpeX on the IRTF.

8 Resolving Power R ~ , SEDs, solid state and molecular bands R ~ 1000 (300 km s-1 ) Atomic features, red shifts and high speed dynamics (e.g., stellar winds) R ~ 104 (30 km s -1 ) line profiles, radial velocities of stars, stellar abundances R ~ 105 (3 km s -1 ) ISM lines, molecular spectroscopy

9 R ~ 1000 CN CaII K CaII H G band Hγ sky (O 2 ) Hβ sky (O 2 ) MgIb Determine the redshifts to galaxies in a cluster with R=500 spectra obtained with LRIS on Keck A B C flux (10 17 erg cm 2 sec 1 Å 1 ) galaxy A z= galaxy B z= Declination (2000) F E 0 galaxy C z= D wavelength (Å) kpc Mullis et al. (2003, ApJ, 594, 154) s 45 s 40 s 35 s Right Ascension (2000) 15 h 24 m 30 s

10 Resolving Power R ~ , SEDs, solid state and molecular bands R ~ 1000 (300 km s-1 ) Atomic features, red shifts and high speed dynamics (e.g., stellar winds) R ~ 104 (30 km s -1 ) line profiles, radial velocities of stars, stellar abundances R ~ 105 (3 km s -1 ) ISM lines, molecular spectroscopy

11 R ~ 10 4 Determine the Be abundances of stars with R=48,000 spectra from HIRES on Keck. Fig. 7. Synthesis fit for Be in KW 227 showing that the Be is present, but depleted. The best fit through the data points is the solid line; the dotted line is for the meteoritic Be abundance, and the dot-dashed line is for essentially no Be. Boesgaard et al. (2004, ApJ, 605, 864)

12 R ~ 10 4 B field strengths of young stars via Zeeman broadening with R=36,000 spectra CSHELL on the IRTF. (Δλ λ 2 ) Fig. 3. Best fit of infrared spectra using a single-component magnetic field model. Magnetically sensitive Ti i lines are fitted where the bold regions indicate the portions of the spectrum actually used in the fit (histogram: data; dash-dotted line: fit without magnetic field; smooth solid line: fit with magnetic broadening). Yang et al. (2005, ApJ, 635, 466)

13 Resolving Power R ~ , SEDs, solid state and molecular bands R ~ 1000 (300 km s-1 ) Atomic features, red shifts and high speed dynamics (e.g., stellar winds) R ~ 104 (30 km s -1 ) line profiles, radial velocities of stars, stellar abundances R ~ 105 (3 km s -1 ) ISM lines, molecular spectroscopy

14 R ~ 10 5 Resolved individual molecular lines.

15 Types of Spectrographs Objective Prism (Long) Slit Spectrograph Multi-Object Spectrograph (MOS) Integral Field (IF) Spectrograph (Fourier Transform Spectrometers) (Heterodyne Spectrometers)

16 Object Prism Spectrograph Used to obtain low resolution spectra of all stars in a field. Image of NGC 2264 in Orion showing stars with Hα (6563 Å) in emission. Courtesy G. Herbig

17 (Long) Slit Spectrograph Work horse of astronomical spectroscopy. Pickering et al. (1999, AJ, 118, 765)

18 Multi-Object Spectrograph Hectospec/MMT Used to obtain a large number of moderate resolution spectra of clusters of stars or galaxies.

19 Integral Field Spectrograph Used to obtained 2D spectra of a small field of view.

20 Spectrograph Basics

21 A Schematic Spectrograph Collimator Camera slit Disperser focal plane Objective Detector at focal plane not to scale

22 A Schematic Spectrograph Collimator Camera 20 μm slit Disperser focal plane Objective Detector at focal plane not to scale

23 A Schematic Spectrograph 0.5 m Collimator Camera slit Disperser focal plane Objective Detector at focal plane not to scale

24 A Schematic Spectrograph 5 cm Collimator Camera slit Disperser focal plane Objective Detector at focal plane not to scale

25 Dispersers Prism - Typically low resolution Grating - Low to very high resolution Grism - (prism+grating) Low to moderate resolution (Fabry-Perot, Michelson Interferometer) (Heterodyne)

26 Diffraction Gratings (a) grating normal + ray 1 ray 2 grating normal + incident light β 1 diffracted light α reflected light β 0 diffracted light A α β β 1 B d d sin β d sin α d Sign Conventions m = d(sin + sin ) Grating Equation

27 Diffraction Gratings m nm m = d(sin + sin ) For a given α and β, the equation holds for many values of m and λ.

28 Diffraction Gratings grating

29 Angular Dispersion ray 1 ray 2 grating normal + m = d(sin + sin ) d sin β A α β B d sin α d angular dispersion A d d = m d cos = sin +sin cos

30 Angular Dispersion A d d = m d cos = sin +sin cos Larger A?: higher order small facets larger angles

31 Buzz Words A d d = m d cos = sin +sin cos Small facets are hard. But, Echelette gratings have small d (1200 lines mm -1 ), work at low m (~1) and small β. Echelle gratings have big d (79 lines mm -1 ), work at high m (~40) and large β.

32 Diffraction Gratings The higher the order, the less wavelength coverage due to overlap.

33 Diffraction Gratings second disperser The higher the order, the less wavelength coverage due to overlap.

34 Cross Dispersion Echelle format

35 Other Dispersers (Prism) low resolution (R~10-500) High efficiency since light is not diffracted into many orders. Used in high efficiency low resolution spectrographs and cross dispersion.

36 Other Dispersers (Grism) Combination of a prism and transmission grating giving R ~ Can be used on axis in a camera or as a cross disperser.

37 Spectrograph Theory

38 Spectrograph Basics What is the R of this system? Collimator Camera w w slit w Disperser (A) focal plane Objective Detector not to scale

39 What is R? How do you want to proceed? 6 slide derivation the answer

40 Δλ w Δl What is R? l = f cam d d = l f cam A = f cam A β fcam The limit of resolution, δλ, is the wavelength difference when the separation (Δl) of the two images is equal to w. = w f cam A R / = f cama w

41 What is w? D w Collimator dcol Camera w slit focal plane Objective ftel fcol fcam What is w in terms of things we know?

42 Quick Review focal plane α α w fcam w = f cam

43 Reimaging w D w Collimator dcol Camera w slit focal plane Objective ftel fcol fcam = w/f cal w = w ( f cam f col )

44 Other useful relations D Collimator Camera w Φ dcol w slit focal plane ftel fcol fcam Objective D f tel = d col f col = w f tel

45 What is w? D w Collimator dcol Camera w Objective ftel slit fcol w = D f cam d col fcam focal plane

46 A bit of math... R / = f cama w w = D f cam d col R / = d cola D

47 Summary angular slit size D Collimator Camera w Φ dcol w slit Disperser (A) focal plane ftel fcol fcam Objective R / = d cola D

48 Resolving Power R / = d cola D For a fixed R, a larger telescope means larger collimator means bigger instrument. Cost scales as the volume. To change R for a given telescope, change A or Φ. 2A 0.5Φ

49 Spectral Data Reduction

50 Spectral Data Reduction Data reduction is, in some sense, more art than science, passed on from advisor to student (hopefully well). There is no correct way to reduce data and how you reduce it often depends on the science you want to do. Nevertheless...

51 Data Reduction How do we go from this: to this?

52 Raw Data What is in the raw data? Spatial (x) Crack OH sky lines Object Bad Pixels Spectral (λ) Ghost

53 Background Subtraction

54 Extraction f i = j D[i, j]

55 Extraction

56 Wavelength Calibration argon emission lines

57 Wavelength Calibration Flux Pixels

58 Wavelength Calibrated Raw Spectrum

59 But what do we actually have? O( ) = [I( ) T ( )] P ( ) Q( ) O(λ) - Observed spectrum I(λ) - Intrinsic spectrum P(λ) - Instrument profile T(λ) - Atmospheric transmission Q(λ) - Instrument throughput

60 I(λ) T(λ) I(λ) T(λ) P(λ) [I(λ) T(λ)] P(λ) [I(λ) T(λ)] P(λ) Q(λ) [I(λ) T(λ)] P(λ) Q(λ) Wavelength (µm) Q(λ) III(λ) Wavelength (µm)

61 How do we find I(λ)? O( ) = [I( ) T ( )] P ( ) Q( ) O( ) = [I( ) P ( )] [T ( ) P ( )] Q( ) O( ) = I instr ( ) R( ) Typically happy with Iinstr(λ), so we just need R(λ).

62 How do we find R(λ)? Observe a standard star. O std ( ) = I std instr ( ) Rstd ( ) If we know I std instr(λ), then we have R(λ) as, R std ( ) = Ostd ( ) I std instr ( )

63 How do we find R(λ)? R std ( ) = Ostd ( ) I std instr ( ) Q(λ) R std (λ) T(λ)

64 Final Spectrum Iinstr sci ( ) = Osci ( ) R std ( ) = Isci instr ( )Rsci ( ) R std ( ) HD (K4 V)

65 Practical Considerations Flux and Telluric Calibrations Parallactic angle Standard stars and wavelength calibration references

66 Flux and/or Telluric Correction? Optical Red Optical Infrared Infrared

67 Flux and/or Telluric Correction? Typically no correction for telluric absorption is done in the optical. At infrared wavelengths, telluric correction is crucial!

68 Parallactic Angle As light passes through our atmosphere, it is refracted. z The amount of refraction depends on the wavelength, zenith angle (z), and the atmospheric conditions. venus

69 Parallactic Angle As light passes through our atmosphere, it is refracted. The amount of refraction depends on the zenith angle (z) and is is along azimuth lines. Dispersion (arcsec) Pressure = 615 mb Precipitable water vapor = 2 mm Atmospheric temperature = 260 K Altitude = 4200 m (MKO) Wavelength (µm)

70 As light passes through our atmosphere, it is refracted. Parallactic Angle seeing The amount of refraction depends on the wavelength, zenith angle (z), and the atmospheric conditions. refraction Must align the slit with the dispersion. The angle changes for equatorial telescopes.

71 zenith NCP E N S W

72 Standard Stars Optical Flux Standards: Massey et al. (1988, ApJ, 328, 315) Massey & Cornwall (1990, ApJ, 358, 344) Hamuy et al. (1994, PASP, 106, 566) A0 V and G2 V stars for near-infrared find_a0v.cgi See Vacca et al. (2003, PASP, 115, 389)

73 Wavelength Calibration Optical Wavelength Calibration He/Ne/Ar emission lines, OH Emission lines, Osterbrock et al. (1998, PASP, 110, 1499) Osterbrock et al. (1997, PASP, 109, 6140) Osterbrock et al. (1996, PASP, 108, 1051) Near-Infrared Calibration He/Ne/Ar emission lines Hinkle et al. (2001, PASP, 113, 548) OH lines, Rousselot et al. (2000, A&A, 354, 1134)