Galaxies at high redshift

The investigation of the metallicity evolution of galaxies is currently limited to z < 4, through the detection of optical emission lines in star-forming galaxies (e.g. Erb et al. 2006; Pettini et al. 2001) or of absorption lines in DLA (Prochaska et al. 2003; Kulkarni et al. 2005). At higher redshift (z ~ 6) the metallicity has been inferred only for a few, exceptionally luminous QSOs (Pentericci et al. 2002), while for the general population of (normal) galaxies the metallicity evolution between z ~ 4 and the epoch of re-ionization (z ~ 7-12) is largely unknown. This crucial redshift range is essentially unexplored because of both low sensitivity and limited spectral coverage of ground-based instrumentation. For instance, the classical R23 parameter (defined as R23 = ([O ii]3727 + [O iii]4959 + [O m]5007)/H|3), which is widely used to trace gaseous metallicities, is observable from the ground only up to z ~ 3.5 (i.e. as long as [O iii]5007 is observable in the K band). JWST, and NIRSpec in particular, will be able to extend these diagnostics to higher redshift and to fainter galaxy populations than is feasible from the ground. As an example, Figure 23.4 shows the simulated NIRSpec spectrum of a low-metallicity galaxy (Z ~ 0.1 ZQ) at z = 8, with mAB = 27, observed with the MSA configuration at R = 1,000 and with an integration time of 105 s: all of the lines required to determine the R23 parameter are clearly detected, therefore providing an estimate of the metallicity. At z > 9 the [O iii]5007 line leaves the NIRSpec range and will be observable only with MIRI, but only for brighter targets, due to the lower sensitivity of the latter. An alternative diagnostic for metallicity that is usable at high redshift (especially for metal-poor galaxies) is the ratio [Ne iii]3869/[O ii]3727. Indeed, the latter line ratio is seen to anti-correlate with metallicity, and at Z < 0.1ZQ the strength of [Ne iii]3869 is comparable to that of [O ii]3727 (Nagao et al. 2006a). These lines are observable in the NIRSpec band up to z ~ 12. The availability of this and other gas-metallicity diagnostics as a function of redshift, both for ground-based instruments and for JWST, is summarized in Figure 23.5.

Li 15

[o iii]4959

[o ii]3727

[Ne iii]3869

i/lfVVAAJVVV,

[Ne iii]3869

l i/lfVVAAJVVV,

Figure 23.4. A simulated NIRSpec spectrum of an mAB = 27 galaxy at z = 8, with a metallicity of Z ~ 0.1Z0, observed with the MSA configuration at R = 1,000 and with an integration time of 105 s. The main, detected emission lines are labeled.

In principle JWST could identify spectroscopically the first generation of stars (Population III), free of metals, through the detection of the He ii emission line at Arest = 1,640 A, which is expected to have an equivalent width of a few tens of ngstrom units in extremely metal-poor systems (Z < 10-4ZSun) (Schaerer 2003). However, as discussed in Panagia (2004), the expected faintness of the primordial galaxies will probably limit the capability to detect the He ii line, except for the bright end of the population (or for lensed objects). The latter may be the case for a few strong Lya emitters at z > 6 detected in wide-area surveys, and characterized by a large EW(Lya), which is suggestive of very young and extremely metal-poor galaxies (Nagao et al. 2006b). Followup observations of these objects have so far failed to detect the He ii line, but with loose constraints due to the limited sensitivity of ground-based observations (Nagao etal. 2005). The JWST-NIRSpec combination will provide much deeper spectra of these objects and, it is hoped, detect the He ii line or, alternatively, detect metal lines that would disprove their Population-III/metal-free nature.

Note that, insofar as normal/star-forming galaxies at high redshift are concerned, JWST will be able to measure global metallicities, while relative elemental abundances will likely remain unconstrained. This is because the observable, strong emission lines at z > 4 do not provide additional information beside the global metallicity (especially at z > 6.6, where even the faint [N ii] is outside the NIRSpec range). For a few exceptionally bright (e.g. lensed) sources it will be possible

Figure 23.5. Availability of some of the gas-metallicity diagnostics as a function of redshift, both for ground-based facilities and for JWST.

to determine abundance patterns from the detection of stellar and ISM absorption features (Pettini et al. 2002; Rix et al. 2004; Mehlert et al. 2006), but these will be relatively rare cases.

Quasi-stellar objects (QSOs) are an alternative class of objects that can provide important information on the relative abundance of elements at high redshift. Even if high-z QSOs are rare, and probing overdense regions, their near-IR spectra (UV-rest frame) are rich in emission and absorption features whose relative strength provides crucial information on the star-formation history in the early Universe. This is especially true for the prominent iron emission features, relative to emission lines of a-elements, which provide a measure of the Fe/a ratio. The latter gives important constraints on the epoch of initial star formation, since a-elements are mostly produced by Type-II SNe, whereas iron is mostly produced by SN Ia, which

Figure 23.6. A simulated JWST-NIRSpec R = 1,000 (continuum-subtracted) spectrum of a QSO at z = 6.5, with a brightness comparable to that of the faintest QSOs found in the SDSS at z > 6 (mAB = 21), for an integration time of 20 min.

Figure 23.6. A simulated JWST-NIRSpec R = 1,000 (continuum-subtracted) spectrum of a QSO at z = 6.5, with a brightness comparable to that of the faintest QSOs found in the SDSS at z > 6 (mAB = 21), for an integration time of 20 min.

provide a delayed enrichment relative to the former; hence the Fe/a ratio is regarded as a sensitive "clock" of star formation. Various authors have tried to investigate the Fe/a ratio in high-z QSOs (e.g. Dietrich et al. 2003; Maiolino et al. 2003; Iwamuro et al. 2005). However, at z > 4 ground-based observations can access only the UV iron emission blend at 2,700-3,050 A, which is not a good tracer of iron abundance, yielding inconclusive results. The optical iron emission at 4,200 A < Arest < 5,400 A (the "optical Fe hump") is a better indicator of iron abundance (Verner et al. 2004; Verner & Peterson 2004), but it is shifted outside of the K band at z > 4, and therefore unobservable at the interesting redshifts approaching re-ionization. JWST will easily detect and measure the intensity of the optical Fe emission in the spectra of QSOs at z > 6, as illustrated in Figure 23.6, thus providing crucial information on the epoch of star formation in these objects.

References

Boisson, C., Joly, M., Pelat, D., & Ward, M. J. (2004), A&A 428, 373

Brandl, B. R., Devost, D., Higdon, S. J. U. et al. (2004), ApJS 154, 188

Bressan, A., Panuzzo, P., Buson, L. et al. (2006), ApJ 369, L55

Dasyra, K. M., Tacconi, L. J., Davies, R. I. et al. (2006), ApJ 651, 835

Dietrich, M., Hamann, F., Appenzeller, I., & Vestergaard, M. (2003), ApJ 596, 817

Egami, E., Dole, H., Huang, J.-S. et al. (2004), ApJS 154, 130

Erb, D. K., Shapley, A. E., Pettini, M. et al. (2006), ApJ 644, 813

Frayer, E., Chapman, S. C., Yan, L. et al. (2004), ApJS 154, 137

Genzel, R., & Cesarsky, C. J. (2000), ARA&A 38, 761

Gardner, J. P., Mather, J. C., Clampin, M. et al. (2006), Space Sci. Rev. 123, 485 Helmi, A., Irwin, M. J., Tolstoy, E. et al. (2006), ApJ 651, L121 Horner, S. D., & Rieke, M. J. (2004), Proc. SPIE 5487, 628-634

Imanishi, M., Dudley, C. C., & Maloney, P. R. (2006), ApJ 637, 114 Iwamuro, F., Kimura, M., Eto, S. et al. (2004), ApJ 614, 69 Kulkarni, V. P., Fall, S. M., Lauroesch, J. T. et al. (2005), ApJ 618, 68 Larsen, S. S., Origlia, L., Brodie, J. P., & Gallagher, J. S. (2006), ApJ 634, 1293 Liang, S., Hammer, F., & Flores, H. (2006), A&A 447, 113

Maiolino, R., Juarez, Y., Mujica, R., Nagar, N. M., & Oliva, E. (2003), ApJ 596, L55 Mehlert, D., Tapken, C., Appenzeller, I., Noll, S., de Mello, D., & Heckman, T. M. (2006), A&A 455, 835

Nagao, T., Motohara, K., Maiolino, R. et al. (2005), ApJ 631, L5

Nagao, T., Maiolino, R., & Marconi, A. (2006), A&A 459, 85

Nagao, T., Murayama, T., Maiolino, R. et al. (2007), A&A 468, 877

Origlia, L., Ferraro, F. R., Fusi Pecci, F., & Oliva, E. (1997), A&A 321, 859

Origlia, L., Goldader, J. D., Leitherer, C., Schaerer, D., & Oliva, E. (1999), ApJ 514, 96

Origlia, L., Valenti, E., Rich, R. M., & Ferraro, F. R. (2006), ApJ 646, 499

Panagia, N. (2004), in IMF50: The Stellar Initial Mass Function Fifty Years Later, eds. E.

Corbelli, F. Palla, & H. Zinnecker (Dordrecht, Springer), pp. 457-479 Pentericci, L., Fan, X., Rix, H.-W. et al. (2002), ApJ 123, 2151 Pettini, M., Shapley, A. E., Steidel, C. C. et al. (2001), ApJ 554, 981 Pettini, M., Rix, S. A., Steidel, C. C. et al. (2002), ApJ 569, 742 Prochaska, J. X., Gawiser, E., Wolfe, A. M., Castro, S., & Djorgovski, S. G. (2003), ApJ 595, L9

Rich, R. M., & Origlia, L. (2005), ApJ 634, 1293

Risaliti, G., Maiolino, R., Marconi, A. et al. (2003), ApJ 595, L17

Rix, S. A., Pettini, M., Leitherer, C. et al. (2004), ApJ 615, 98

Rowlands, N., Evans, C., Greenberg, E., et al. (2004), Proc. SPIE 5487, 676-687

Savaglio, S., Glazebrook, K., Le Borgne, D. et al. (2005), ApJ 625, 260

Shi, Y., Rieke, G. H., Hines, D. C. et al. (2006), ApJ at press (astro-ph/0608645) Smith, J. D. T., Dale, D. A., Armus, L. et al. (2004), ApJS 154, 199 Spoon, H., Armus, L., Cami, J. et al. (2004), ApJS 154, 184 te Plate, M., Holota, W., Posselt, W. et al. (2005), Proc. SPIE 5904, 185 Valenti, E., Ferraro, F. R., & Origlia, L. (2004), MNRAS 351, 1204 Valenti, E., Origlia, L., & Ferraro, F. R. (2005), MNRAS 361, 272 Verner, E., Bruhweiler, F., Verner, D., Johansson, S., Kallman, T., & Gull, T. (2004), ApJ 611, 780

Verner, E., & Peterson, B. A. (2004), ApJ 608, L85

Willner, S. P., & Nelson, K. (2002), ApJ 568, 679

Wright, G. S., Rieke, G. H., Colina, L. et al. (2004), Proc. SPIE 5487, 653

Was this article helpful?

0 0

Post a comment