1. The Low-End MF of the Arches Cluster Sungsoo S. Kim Kyung Hee Univ., Korea with Jihye Shin Don Figer

2. The Arches Cluster ~30 pc away from the Galactic Center (in projection) Young, compact cluster –Age = ~2.5 Myr old –Mass = 3-4x10 4 M ⊙ –Size = ~1 pc (from King-profile fit) One of the most massive & youngest clusters in the Milky Way

3. Star Formation in the GC Harsh environment –Strong tidal forces –Large internal velocity dispersions within clouds –Strong magnetic fields A different mode of star formation is expected in the Galactic Center Large  F g, T, B  External Compression Required  Flatter IMF ? Elevated IMF Lower Limit ?

4. The Arches Cluster as a Benchmark of Infrared Cameras

5. Lick 3-m Figer (1995)

6. Keck I 10-m Serabyn et al. (1998)

7. HST/NICMOS Figer et al. (1999)

8. Gemini Hokupa’a (Demo) Stolte et al. (2002)

9. VLT NAOS/CONICA Stolte et al. (2005)

10. Keck II NIRC2-LGS Kim et al. (2006)

11. Luminosity & Mass Functions from the Most Recent Observation

12. Kim, Figer, Kudritzki, & Najarro (2007) Keck II NIRC2-LGS May 2006 One of the first crowded- field observations with the Keck LGS/AO 0.02”/pix, 1024x1024 900s exp., H & K s Medium & wide cameras 9-point dither pattern 0.06~0.25 Strehl ratio 0.06”~0.10” FWHM

13. Moderate radial elongation of PSF due to a small isoplanatic patch Less extinction in the core – winds or from massive stars? – photo-evaporation by intense UV radiation?

14. One cluster field + three control fields Controls are needed to estimate the background population. No NGS near the cluster  Need for LGS/AO

15. Reduced with DAOPHOT mag error 0.05 down to K s 19 Calibrated against HST/ NICMOS obs. 2Myr solar Z isochrone for the cluster 1Gyr solar Z isochrone for the control

16. Luminosity functions: Controls (Dereddened, completeness-corrected, recovery fraction>50%)

17. Luminosity functions: Cluster & Controls (Dereddened, completeness-corrected)

18. Mass Functions (Keck: Background-subtracted, VLT: Includes background) Γ=-0.9

19. Bump in the NGC 3603 (Eisenhauer 1998)

20. Is the bump real? Eisenhauer et al. (1998) adopted a PMS mass-mag relation by Palla & Stahler (1993), which has a steep slope at 3-7 M ⊙. But, PMS mass-mag relations are thought to have rather large uncertainties with –the extinction by the dust envelope –the assumed accretion rate during the star formation.

21. Evolution of the MF The Arches cluster has dynamically evolved: –High density  short relaxation timescale –Strong tidal field, large stellar mass range  makes it evolve even faster, –Thus even at the age of 2.5 Myr, the PDMF may be quite different from the IMF. We performed Fokker-Planck & Nbody6 simulations: –Circular orbits assumed. –Results are insensitive to the lower mass cutoff (0.1, 1 M ⊙ ). –Power-law slope  decreases by 0.1–0.2 during 2–2.5 Myr. –Thus the IMF slope is estimated to be 2.0–2.1. 1.3–50 M ⊙ ; 5”–9” annulus

22. The Latest MF of the Arches Without the proper control field data, it was previously thought that –there may be a turnover near 6.3 M ⊙, –this indicates a global decrease of the cluster MF. But we find that –the 6.3 M ⊙ turnover is simply a local bump in the MF, –the MF continues to increase at least down to 1.3 M ⊙, –the present-day MF has a power-law slope   1.9 for 5”–9” annulus. The bump was also observed in another young cluster, NGC 3603. Can we get more information below 1.3 M ⊙ ?

23. Theoretical Studies on the Star Formation in the GC (1/2) Klessen et al. (2007, MNRAS, 374, L29) –Hydrodynamic simulations for warm, dusty molecular clouds. –Found a top-heavy IMF with a sharp turn-down below 7 M ⊙.

24. Theoretical Studies on the Star Formation in the GC (2/2) Dib et al. (2007, MNRAS, 381, L40) –Semi-analytic calculations to compare the local rates of collapse and coalescence in the molecular clouds. –Was able to find conditions that result in 6 M ⊙ bump in the pre-stellar core MF.

25. The Pixel Intensity Histogram Method for the Low-End Mass Function

26. The Pixel Intensity Histogram Method Conventional Photometry –Identifies resolved stars and assigns magnitudes to them. –Unable to recover  too faint stars (background-limited)  stars in crowded regions (confusion-limited) Pixel Intensity Histogram Method –Uses the intensity histograms of pixels with low intensities, i.e., the pixels that do not belong to the bright, resolved stars. Pixels that belong to resolved stars

27. –One can find the model mass function that best fits the observed pixel histogram. –One needs to model not only the cluster stars, but also the background and the field stars. –For the Arches cluster, we use the HST/NICMOS data because the characteristics of the images are relatively well known. Other Statistical Methods –Count-Brightness Ratio (Baum & Schwarzschild 1955) Estimates the stellar content (population) of external galaxies. Analyzes the ratio of the number of resolved stars to the surface brightness. –Surface Brightness Fluctuations (Tonry & Schneider 1988) Measures distances of distant galaxies. Analyzes pixel-to-pixel intensity variation. Needs to have a large number of stars falling onto each pixel. –Two-Point Correlation Function (e.g., Peebles 1980) Measures the clustering of galaxies. Analyzes the distribution of resolved objects.

28. Components of the Pixel Intensities Readout Noise –Mean = 31 e -, STDEV = 5.7 e - Dark Currents –0.03 e - /s → 8 e - for 256 sec. Zodiacal Light –~0.2 e - /s → ~40 e - for 256 sec. Thermal Radiation from the Telescope –~100 e - /s → ~25,000 e - for 256 sec. Background and Foreground Stars –Observed control field LF + Star count model Cluster Stars –Observed cluster LF + Various model MFs Diffuse Near-IR Background –Scattered off dust grains? –~5,000 e - for 256 sec.

29. Modeling the Control Fields (1/3) 1.Create a synthetic LF –with the observed LF from the HST and Keck observations for the bright part –with a Milky Way star count model (Wainscoat et al. 1992) for the faint part HST Keck Model Combined

30. Modeling the Control Fields (2/3) 2.Create an artificial image –with model PSFs from TinyTim (80x80 pixel) –with a thermal background level that fits the observation Observation Artificial

31. Modeling the Control Fields (3/3) 3.Add dark clouds to match the lowest intensity pixels –by introducing a certain amount of diffuse IR radiation ObservationArtificial Image Bright-Pixel Masked Images

32. Modeling the Cluster Field (0.2−0.35 pc annulus) Create synthetic LFs –with the observed LF from the HST and Keck observations for the bright part –with several power-law MFs for the faint part HST Keck Power-law Combined

33. α = −1 ( Γ = 2 ) MF LF Pixel Histogram

34. α = 1 ( Γ = 0 ) MF LF Pixel Histogram

35. α = 2 (Γ = −1 ) MF LF Pixel Histogram

36. α = 3 (Γ = −2 ) MF LF Pixel Histogram

37. Polarization toward the Cluster Observed with IRSF/ SIRPOL and reduced by Shogo Nishiyama Field lines are perpendicular to the galactic plane. Ejection of the remnant gas material along the field lines?

38. Proper Motions Observed by Stolte et al. (2007) at VLT and Keck with a 4.3 year baseline (24 mas). The cluster is moving parallel to the Galactic plane. The orbital direction does not coincide with the direction of less extinction. –Thus the magnitude of extinction seems to have nothing to do with the orbital motion. Stolte et al. (astro-ph/0706.4133)

39. The pixel histogram method can be used to estimate the stellar content below the completeness limit of the PSF photometry. The MF in the 0.2-0.35 pc annulus has  = 2−3 at least down to 0.5 M ⊙. The integrated luminosity from stars with M<0.3 M ⊙ is too small that even the pixel histogram method is not applicable. The existence of a relatively large amount of diffuse IR radiation remains to be examined. One could apply this method to other star clusters with unresolved, faint stellar population. Results from the Pixel Histograms

40. From the conventional PSF photometry, we know that the PDMF of the Arches cluster rises at least down to ~1 M ⊙. The Arches cluster is dynamically evolved, but simulations show that its IMF is not much different from the PDMF in the intermediate annulus, which is not much different the Salpeter IMF (  = 2.1 down to 1 M ⊙ ). Analysis of pixel intensity histograms suggests a non-negligible fraction of stars below 1 M ⊙ (  ~ 2−3 down to 0.5 M ⊙ ). Summary