1. APOGEE: The Apache Point Observatory Galactic Evolution Experiment l M. P. Ruffoni 1, J. C. Pickering 1, E. Den Hartog 2, G. Nave 3, J. Lawler 2, C. Allende-Prieto 4 1.Imperial College London, UK 2.University of Wisconsin, Madison, WI, USA 3.NIST, Gaithersburg, MD, USA 4.University of Texas, Austin, TX, USA 1

2. DurationSpring 2011 to Summer 2014 Spectra Measuring 1.51µm <  < 1.7µm Resolving power ~30,000 S/N Ratio greater than 100 Targets100,000 evolved stars 15 elements - Fe most important PrecisionMetal abundances to ~0.1 dex Radial velocities to <0.3 km s -1 2 APOGEE is one of 4 instruments forming the third Sloan Digital Sky Survey (SDSS3) It will conduct a spectroscopic survey of all stellar populations in the Milky Way It will measure in the near-IR where Galactic dust extinction is ~1/6 of that at visible wavelengths It will measure chemical abundances and radial velocities of 100,000 evolved stars to help explain Galactic evolution

3. Detecting elements in stars Photosphere Hot, dense interior Emission contains absorption lines Section of a star Visible spectra for different star types n =  E = 0 n = 4 E = -0.85 eV n = 3 E = -1.51 eV n = 2 E = -3.40 eV n = 1 E = -13.6 eV Type Absorption lines indicate the presence of an element. Line strength is mainly linked to: Stellar properties (e.g. temperature) Absorption transition probability Chemical abundance Temperature/KType K A O Effect of lower level population on H  3

4. Why look at QSOs when studying  ? Determining chemical abundances 1)Use a  2 fit to stellar models to find Stellar temperature Surface gravity Microturbulence parameter Abundance of important elements [Fe/H], [C/H], and [O/H] 2) Fix these parameters then fit other abundances Simulated H-band spectra for an Fe-poor (black) and Fe-rich (grey) star. All other parameters fixed. 4 Experimentally measured transition probabilities in the literature: J. C. Pickering et al. Can J Phys 89 pp. 387 (2011) ScTiVCrMnFeCoNiCuZn –457–2651–4–1

5. Measuring Transition Probabilities E2E1E2E1  A 21 B 12 B 21 Einstein coefficients Spontaneous AbsorptionStimulated Emission Emission Transition probabilities can be obtained from emission spectra 5 To vacuum pump High voltage ~600 V at 50-1000 mA Ne or Ne/Ar gas at 1-3 mbar Glass AnodeCathodeAnode

6. Decay to a single level Decay to multiple levels E2E1E2E1 E2E1E2E1 I I I I 12 Branching Fractions 12 BF = Branching fraction EW = Line equivalent width 6

7. Free spectral range Spectral range determined by Spectrometer optics Detector sensitivity Filter combinations Measurement electronics Complications 1.0 0.0 Spectrometer Response Determined by measuring a calibrated continuum source Tungsten lamp (IR to UV) Deuterium lamp (UV and vacuum UV) I 1.0 0.0 I Normalised response 0 4000 8000 12000 35000 45000 55000 Wavenumber / cm -1 1.0 0.8 0.6 0.4 0.2 0.0 Normalised Response W lamp D 2 Lamp Either Select a range to measure all upper level branches or Use overlapping spectra to carry calibration 7

8. Branching Fractions for APOGEE A 21 for Neutral Fe (Fe I ) is of the greatest importance Wavenumber / cm -1 S/N S/N x Calibration function  / cm -1 EWBF / % 6124.101692413.499537 ± 3 6395.4032910599.5889732 ± 2 6742.8734919624.8279460 ± 2 Target H-band lines in a single spectral range Extract all lines from a single upper level Calibrate line intensities Fit line profiles to get BFs 8

9. Catch-22: Branching fractions or level lifetimes Laser induced fluorescence (LIF) is used to measure  2 E2E1E2E1 205 - 720 nm UV - visible No significantly populated lower level to excite Not possible to measure  2 0 4000 8000 12000 35000 45000 55000 Wavenumber / cm -1 1.0 0.8 0.6 0.4 0.2 0.0 Normalised Response Change target levels so they are LIF compatible No lines to carry calibration Not possible to measure BFs 9

10. Solutions Some levels contain visible lines to link IR spectra to UV spectra LIF has a low lying level available FTS line intensities can be calibrated For FTS compatible upper levels Use theory to calculate  2 (check against known levels) Constrain models with stellar spectra For LIF compatible upper levels Use lines from similar upper levels as calibration proxies Link spectral regions with theory calculations Use stellar spectra to estimate missing calibration factors The work continues... (with the support of the STFC) 10