ABSTRACT. We present high-resolution spectrophotometric observations of the Orion Nebula, made with the. Cassegrain echelle spectrograph on the Blanco 4 ...
THE ASTROPHYSICAL JOURNAL SUPPLEMENT SERIES, 129 : 229È246, 2000 July ( 2000. The American Astronomical Society. All rights reserved. Printed in U.S.A.
HIGH-RESOLUTION SPECTROSCOPY OF FAINT EMISSION LINES IN THE ORION NEBULA J. A. BALDWIN,1 E. M. VERNER,2,3 D. A. VERNER,2 G. J. FERLAND,2 P. G. MARTIN,3 K. T. KORISTA,2,4 AND R. H. RUBIN5 Received 1999 December 21 ; 2000 February 1
ABSTRACT We present high-resolution spectrophotometric observations of the Orion Nebula, made with the Cassegrain echelle spectrograph on the Blanco 4 m telescope at Cerro Tololo Inter-American Observatory (CTIO). The resolution and signal-to-noise ratio make it possible to identify 444 emission lines in the 3498È7468 A range, down to 104 times fainter than Hb. We present a detailed atlas of these emission lines along with an analysis of the associated errors. This data set is used to study the velocity Ðeld in the Orion Nebula. The forbidden lines split into two distinct groups. The low-ionization group has ions with an ionization potential less than 20 eV. Lines of these ions, [O I], [N I], [Ni II], and [Fe II], have recession velocities, relative to the hydrogen lines, of ]10 to ]15 km s~1. There is a sharp change to the second, high-ionization group, which includes lines of ions with ionization potentials larger than 20 eV, namely, [S II], [O II], [N II], and [Fe III]. These lines have velocities around ]3 km s~1, with a slight trend of decreasing velocity with the increasing ionization potential. This is consistent with previously proposed dynamical models in which lines of ions with di†erent ionization potentials originate at di†erent distances from the ionizing stars. SigniÐcant acceleration appears to take place across the narrow region where Fe2` exists. Across this region the gas receives an acceleration of D 2.5 ] 10~5 cm s~2. This provides a constraint on hydrodynamical models. We set a limit He II 4686/Hb \ 7 ] 10~5, which in turn sets a limit to the intensity of the ionizing continuum at energies higher than 54 eV. Modern stellar atmospheres predict a continuum that is far stronger than is present in the region near h1 Ori C. Subject headings : H II regions È ISM : individual (Orion Nebula) È ISM : kinematics and dynamics È line : identiÐcation 1.
INTRODUCTION
Pradhan, & Osterbrock 1994), which would be hard to understand if the gas is in photoionization equilibrium. We present new high-resolution and high signal-to-noise observations for almost the same wavelength range as Esteban et al. (1998) but for the position studied and modeled by BFM. Our observations increase the number of measured lines by a factor of 2. [O I] j 5577 was detected with a very low I(5577)/I(6300 ] 6363) ratio, as was predicted by Baldwin et al. (1996, hereafter B96). The main goal of this paper is to make this data set and our line identiÐcations available in the literature for general use. In addition, the high accuracy of the wavelength measurements and detection of many lines allows us to present a detailed picture of the velocity Ðeld. The velocity Ðeld can help identify whether a line forms by collisional excitation or by recombination since di†erent parent species are involved. Our [Fe II] emission-line identiÐcations are based on predictions from the detailed quantitative modeling (B96). The analysis of [Fe II] emission and of other aspects of this data set will be the subject of subsequent papers (Verner et al., in preparation).
The Orion Nebula is perhaps the best-observed Galactic H II region, and extensive line identiÐcation work has been done (Kaler, Aller, & Bowen 1965 ; Peimbert & TorresPeimbert 1977 ; Baldwin et al. 1991, hereafter BFM ; Osterbrock, Tran, & Veilleux 1992, hereafter OTV). The number of identiÐed lines in the optical range has increased dramatically in the last decade. Recently, Esteban et al. (1998) have measured the intensities of about 220 emission lines in the 3500 to 7060 A range, about twice more than OTV. Esteban et al. (1999) used long-exposure CCD high spectral resolution spectrograms to obtain an accurate measurement of the nebular [O I] j 5577 emission line, which is an important diagnostic indicator of physical conditions where the [O I], [N I], and [Fe II] lines form. They found that these lines are produced in the gas with electron density 2000 \ n \ 40,000 cm~3 and temperature 8900 \ T \ e 12,400 K.e This is important because it refutes the claim that the [Fe II] lines form in a dense warm region (Bautista,
2.
1 Cerro Tololo Inter-American Observatory, National Optical Astronomy Observatories, Casilla 603, La Serena, Chile. NOAO is operated by AURA, Inc., under contract to the National Science Foundation. 2 Department of Physics and Astronomy, University of Kentucky, Lexington, KY 40506-0055 ; katya=pa.uky.edu. 3 Department of Astronomy, and Canadian Institute for Theoretical Astrophysics, University of Toronto, Toronto, ONT, Canada M5S 3H8. 4 Western Michigan University, Department of Physics, Kalamazoo, MI 49008. 5 NASA/Ames Research Center, Astrophysics Branch MS 245-6, Mo†ett Field, CA 94035-1000.
OBSERVATIONS AND REDUCTIONS
We obtained two deep echelle spectra (red and blue) on two di†erent nights approximately one year apart. We used the Cassegrain echelle spectrograph on the Blanco 4 m telescope at Cerro Tololo Inter-American Observatory (CTIO). The spectrograph slit was placed at a \ 05h35m13s. 9, 2000 d \ [05¡23@23@@ by Ðrst centering h1 Ori C in the slit 2000then o†setting 37A west. This is the same position and studied by B96 with earlier echelle spectroscopy but is 4A 229
230
BALDWIN ET AL.
north, 1A west of position 2 along the slit location studied extensively at low resolution by BFM.6 Our red spectrum was obtained on UT 27 Feb 1997 using a 31.6 grooves mm~1 echelle grating and a 316 grooves mm~1 cross-disperser (grating G181) with a GG495 order separator Ðlter. This gave full spectral coverage over the wavelength range 5100È7485 A . The slit width was 1A. 0, giving a nominal resolution at the center of each echelle on the order of 10 km s~1. The decker length was 13A. 1, but the area actually extracted along the slit was 1@@ ] 10@@. The slit was oriented along the parallactic angle, P.A. 142¡. A series of exposures, of lengths 10 s, 100 s, and Ðnally three times by 1000 s, were taken so that measurements of strong emission lines from the shorter exposures could be combined with those of weak emission lines from the sum of the three 1000 s exposures without having saturation problems. In practice, only the Ha and the [N II] lines had to be measured from the 10 s exposure, while all other lines were unsaturated on the 1000 s exposures. The blue spectrum was taken on 1998 January 17 using a 79 grooves mm~1 echelle grating and 316 grooves mm~1 cross-disperser (grating KPGL2) with no order separator Ðlter. This spectrum completely covered the wavelength range 3510È5940 A . The slit center was placed at the same position as for the red spectrum, and the slit again rotated to the parallactic angle, which in this case was P.A. 203¡. Again we used a 1A. 0 slit width and extracted data from 10A along the slit. Eight exposures were taken of the Orion Nebula, with lengths of 100 s, six times by 1000 s, and 100 s again. Orion crossed the meridian halfway through this sequence. On both nights, we also took two 500 s exposures at a sky position 4m in h.a. east of the Orion position. The sky spectrum was not subtracted from the Orion spectrum but, rather, was used as a reference in the identiÐcation of sky lines in the Orion data. On each night we also observed two additional positions in the Orion Nebula, which have not yet been analyzed. The spectra were reduced using the IRAF echelle package. A wavelength alignment was then made between the red and blue spectra by comparing the wavelengths of 28 lines measurable in both spectra in the jj5100È5940 region where they overlapped. We expected a di†erence of 10.9 km s~1 caused by the EarthÏs orbital motion ; the actual adjustment needed was j [ j \ 10.3 km s~1. A red blue using the median Ðducial zero velocity was then determined of the velocities of the Ðrst six Balmer lines and then subtracted from each measured velocity. This corresponds to a heliocentric velocity of ]11.9 km s~1. Comparison of our velocities for a number of strong lines of other ions with previous measurements shows agreement within 1 km s~1 (C. R. OÏDell 1999, private communication). For example, we Ðnd ]12.8 ^ 0.6 km s~1 for the [O III] j5007 and j4959 lines, while near our slit position Castan8 edaÏs (1988) measurements for the [O III] j5007 sharp line components (his Table 15) suggest ]12 ^ 0.5 km s~1 for our lower resolution. The Ñux calibration for each night was performed using spectra of the standard stars g Hya, h Crt, and 108 Vir, 6 The positions given in column (1) of Table 2 of BFM should have an o†set of 9A east, 4A south added to them. This is caused by an error in the position of the Ðducial star used by BFM, since clariÐed by C. R. OÏDell (1999, private communication).
FIG. 1.ÈSample of the fully calibrated echelle spectrum in the Balmer limit region illustrating the high resolution and signal-to-noise ratio. Hydrogen lines up to n \ 28 are detectable.
FIG. 2.ÈSample of the fully calibrated echelle spectrum coinciding in wavelength with Fig. 3 of Esteban et al. (1998) and showing the individual emission lines of multiplet 1 of O II as well as much weaker N III, N II lines. A few sharp cosmic-ray hits are present too.
taken through a 10A wide slit oriented in the parallactic angle, and comparing them to the low-resolution calibrating spectra taken over contiguous 16 A intervals by Hamuy et al. (1994). The calibrated one-dimensional spectra for the 1000 s exposures were then co-added using median Ðltering to reject cosmic rays. Samples of the fully calibrated spectrum are presented in Figures 1 and 2. These are purposely chosen to coincide in wavelength with similar plots shown by Esteban et al. (1998) so that readers can easily compare the resolution and signal-to-noise ratios (S/N) in those two sets of echelle data for Orion. We used Gaussian Ðts to measure the peak line wavelength and full width at half-maximum intensity (FWHM). However, it can be seen for the brighter lines that the line proÐles are clearly non-Gaussian, so the integrated intensities for all lines were measured by adding the Ñux above the continuum level within the wavelength interval where the lines were clearly above the continuum noise level.
TABLE 1 EMISSION-LINE MEASUREMENTS FROM 3498 TO 7468 A Rest Wavelength (A ) (1) 3498.660 3512.496 3530.468 3530.495 3554.416 3587.274 3599.273 3613.636 3634.251 3653.554 3653.880 3654.285 3654.679 3655.121 3655.600 3656.114 3656.674 3657.278 3657.936 3658.650 3659.435 3660.291 3661.231 3662.258 3663.408 3664.693 3666.098 3667.685 3669.465 3671.476 3673.759 3676.361 3679.350 3682.803 3686.819 3689.050 3691.558 3694.173 3697.158 3703.850 3705.004 3711.971 3713.049 3717.738 3721.839 3726.062 3728.813 3732.877 3734.362 3735.485 3749.407 3750.144 3764.277 3768.761 3770.627 3777.208 3781.988 3784.856 3789.213 3797.894 3805.738 3806.487 3818.744
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
S/N Ratio (9)
3498.64 3512.51 ... 3530.49 3554.42 3587.27 ... 3613.642 3634.24 3653.608 3653.96 3654.346 3654.756 3655.12 3655.59 3656.11 3656.66 3657.27 3657.92 3658.64 3659.42 3660.28 3661.22 3662.26 3663.40 3664.68 3666.10 3667.68 3669.46 3671.48 3673.76 3676.37 3679.36 3682.81 3686.83 ... 3691.56 3694.212 3697.15 3703.86 3705.00 3711.97 3713.080 3717.717 3721.95 3726.032 3728.815 3732.86 3734.37 ... 3749.484 3750.15 ... ... 3770.63 3777.134 3781.942 3784.89 3789.176 3797.90 3805.765 3806.526 ...
He I He I ... He I He I He I ... He I He I HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI ... HI Ne II HI HI He I HI Ne II S III HI [O II] [O II] He I HI ... O II HI ... ... HI Ne II Fe I He I Fe I] HI He I Si III ...
... ... ? ... ... ... ? ... ... ? ? ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ? ... ... ... ... ... ... ? ... ... ... ... ... ... ? ... ... ? ... ... ? ? ... ? ... ... ? ?
1.7 [1.2 ... 0.4 [0.4 0.3 ... [0.5 0.9 [4.5 [6.6 [5.1 [6.3 0.1 0.9 0.4 1.2 0.7 1.3 0.8 1.2 0.9 0.9 [0.2 0.6 1.0 [0.1 0.4 0.4 [0.3 [0.1 [0.8 [0.8 [0.6 [0.9 ... [0.2 [3.2 0.6 [0.8 0.4 0.1 [2.5 1.7 [8.9 2.4 [0.2 1.4 [0.6 ... [6.1 [0.5 ... ... [0.2 5.9 3.6 [2.7 3.0 [0.4 [2.1 [3.1 ...
23 20 14 19 19 21 12 17 21 18 18 16 17 19 21 20 22 22 24 25 25 25 25 27 25 28 25 26 26 26 26 26 26 26 28 23 27 18 27 28 20 26 22 15 36 15 16 21 25 60 21 26 18 22 26 29 32 19 26 25 17 16 13
0.0233 0.0265 0.0239 0.0290 0.0447 0.0657 0.0033 0.0861 0.1028 0.0052 0.0055 0.0061 0.0094 0.0114 0.0152 0.0191 0.0255 0.0286 0.0344 0.0350 0.0377 0.0452 0.0504 0.0596 0.0620 0.0787 0.0789 0.0881 0.0997 0.1051 0.1174 0.1304 0.1479 0.1636 0.1792 0.0111 0.2233 0.0097 0.2440 0.2845 0.1401 0.3037 0.0130 0.0069 0.5220 10.3841 4.8506 0.0081 0.3898 0.0134 0.0168 0.6863 0.0072 0.0045 0.7653 0.0025 0.0031 0.0060 0.0022 0.8080 0.0061 0.0035 0.0014
0.0417 0.0475 0.0428 0.0519 0.0800 0.1172 0.0059 0.1533 0.1825 0.0092 0.0098 0.0108 0.0167 0.0202 0.0270 0.0339 0.0453 0.0508 0.0610 0.0621 0.0669 0.0802 0.0894 0.1057 0.1099 0.1395 0.1398 0.1561 0.1767 0.1862 0.2080 0.2308 0.2618 0.2895 0.3171 0.0196 0.3945 0.0171 0.4310 0.5016 0.2470 0.5352 0.0229 0.0122 0.9196 18.2900 8.5426 0.0142 0.6852 0.0236 0.0295 1.2047 0.0126 0.0079 1.3400 0.0044 0.0054 0.0105 0.0038 1.4087 0.0106 0.0061 0.0024
7.7 21.4 12.5 12.4 37.3 34.0 4.8 94.4 58.3 6.7 7.1 8.4 12.9 14.2 18.6 24.2 31.7 35.0 39.0 40.1 44.4 53.1 60.7 69.5 79.7 95.9 103.8 122.3 141.4 164.1 176.7 177.0 177.8 190.8 120.1 11.6 129.8 9.9 143.6 156.5 106.1 184.7 8.1 12.6 488.4 19282.3 8998.5 12.5 523.7 8.7 12.3 486.2 6.5 9.6 593.2 4.2 4.9 15.0 4.5 1701.8 17.5 10.3 5.6
Notes (10)
average average
average
average
average blend average identiÐed by Grandi 1976 average average average average average identiÐed by Grandi 1976 large FWHM
average ; 2 orders in poor agreement average
average from the same multiplet as 3694.212 large FWHM
TABLE 1ÈContinued Rest Wavelength (A ) (1)
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
S/N Ratio (9)
3819.620 . . . 3831.697 . . .
3819.61 3831.815 3831.726 ... 3835.384 3837.726 3838.374 3853.665 3856.018 3860.619 3862.595 3867.48 3868.75 3871.79 3878.18 3882.194 3884.597 3889.049 3888.70 3918.968 3920.681 3926.544 3928.556 3935.945 ... 3964.729 3967.460 3970.072 3973.256 3983.723 3985.924 3993.059 3998.63 3998.759 4008.35 4009.257 4023.98 4026.19 ... 4068.600 4069.882 4072.153 4075.862 4076.35 4078.842 4079.70 4083.899 4085.112 4087.153 4089.288 4089.295 4092.929 4095.644 4096.526 4097.225 4097.257 4101.71 4104.990 4110.786 4114.470 4116.104 4119.217 4120.278 4120.816
He I S III C II ... HI S III N II Si II Si II S III Si II He I [Ne III] He I He I O II Co I HI He I C II C II He I S III He I ... He I [Ne III] HI O II S III S III [Ni II] N III S II [Fe III] He I He I He I ... [S II] O II O II O II [S II] O II [Fe III] O II O II O II O II O III O II O II O II O II O II HI O II O II [Fe II] Si IV O II O II He I
... ... ... ? ... ? ... ... ... ? ... ... ... ... ? ? ? ... ... ... ... ... ... ... ? ... ... ... ? ... ... ? ? ? ... ... ... ... ? ... ... ... ... ... ... ... ... ... ? ... ... ... ? ? ... ... ... ... ... ? ? ... ... ...
0.7 ... ... ... [0.2 0.2 [5.2 2.4 1.8 [0.2 1.1 1.1 [0.5 [0.5 [0.1 0.2 1.0 ... ... [3.6 [3.4 [0.3 [0.3 [0.2 ... [0.3 [1.0 [0.2 [0.9 0.0 0.7 15.5 ... ... 0.5 [0.1 [0.3 0.8 ... 4.6 [6.0 [0.4 [0.1 4.8 [0.8 [0.6 [0.7 [0.4 [0.6 ... ... [1.0 1.4 0.7 ... ... 2.0 [3.2 [2.3 7.0 [6.1 [1.5 0.1 0.2
19 32 ... 18 26 17 22 24 19 14 18 18 13 16 15 12 16 35 ... 15 15 17 14 18 17 17 13 25 11 14 11 10 13 ... 16 17 15 18 36 18 34 14 13 18 14 16 15 16 14 14 ... 13 14 13 15 ... 26 19 18 20 15 14 24 19
0.1863 0.0088 ... 0.0111 1.5007 0.0043 0.0095 0.0036 0.0266 0.0020 0.0148 0.0151 2.6885 0.0126 0.0019 0.0024 0.0027 2.9943 ... 0.0094 0.0174 0.0218 0.0030 0.0023 0.0024 0.1920 1.3926 4.5235 0.0013 0.0031 0.0022 0.0033 0.0013 ... 0.0060 0.0330 0.0045 0.4068 0.0042 0.2617 0.0185 0.0115 0.0113 0.0838 0.0019 0.0017 0.0017 0.0027 0.0013 0.0044 ... 0.0017 0.0014 0.0017 0.0065 ... 5.8677 0.0030 0.0037 0.0015 0.0012 0.0061 0.0023 0.0361
0.3238 0.0153 ... 0.0192 2.6011 0.0075 0.0165 0.0062 0.0460 0.0035 0.0255 0.0260 4.6349 0.0217 0.0033 0.0041 0.0046 5.1471 ... 0.0160 0.0297 0.0372 0.0051 0.0039 0.0041 0.3252 2.3576 7.6475 0.0022 0.0052 0.0037 0.0056 0.0022 ... 0.0101 0.0554 0.0075 0.6803 0.0070 0.4338 0.0307 0.0191 0.0187 0.1387 0.0031 0.0028 0.0028 0.0045 0.0021 0.0073 ... 0.0028 0.0023 0.0028 0.0107 ... 9.6639 0.0049 0.0061 0.0025 0.0020 0.0100 0.0038 0.0592
290.6 12.6 ... 23.6 1424.9 9.1 16.1 7.7 67.8 6.6 40.9 41.3 9407.0 38.0 6.3 8.7 6.9 2223.2 ... 38.3 72.6 84.8 13.2 7.6 6.1 324.3 1702.0 2592.5 5.5 11.6 10.3 18.6 7.1 ... 31.0 168.2 16.3 855.8 7.2 867.8 38.2 50.0 56.9 329.5 9.2 7.9 8.0 12.4 6.2 20.4 ... 7.8 5.7 6.6 23.8 ... 16801.0 7.8 17.1 6.6 6.7 34.5 8.9 163.1
3833.544 3835.381 3837.729 3838.307 3853.696 3856.041 3860.617 3862.610 3867.494 3868.743 3871.784 3878.178 3882.196 3884.610 3889.017
... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
3918.921 3920.636 3926.540 3928.553 3935.943 3952.712 3964.724 3967.447 3970.069 3973.245 3983.724 3985.933 3993.265 3998.734
... ... ... ... ... ... ... ... ... ... ... ... ... ...
4008.356 4009.255 4023.976 4026.201 4027.133 4068.663 4069.801 4072.148 4075.860 4076.415 4078.832 4079.692 4083.890 4085.106 4087.145 4089.295
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
4092.915 4095.662 4096.535 4097.265
... ... ... ...
4101.738 4104.947 4110.754 4114.566 4116.020 4119.196 4120.279 4120.818
... ... ... ... ... ... ... ...
232
Notes (10) average large FWHM gh average gh pumped line ; possible blend with S III 3838.31 blend with pumped N II line 3856.06
gh
strongest line from the multiplet average ; large FWHM ; blend pumped line ? pumped line ?
average
average broad bump on He I wing, after subtracting ghost blend with O II 4069.90 line
average gh ; possible blend with O II 4105.00 strongest line in multiplet
TABLE 1ÈContinued Rest Wavelength (A ) (1) 4121.407 4127.880 4128.575 4129.253 4130.896
... ... ... ... ...
4131.832 4132.775 4143.754 4146.052 4153.277 4156.319 4169.018 4177.405 4179.218 4185.439 4189.765
... ... ... ... ... ... ... ... ... ... ...
4201.376 4205.581 4211.289 4219.747 4236.980
... ... ... ... ...
4241.395 4241.762 4244.145 4249.039 4253.539
... ... ... ... ...
4267.168 . . .
4275.554 . . . 4276.987 . . . 4284.931 . . . 4287.587 . . . 4294.339 . . . 4294.777 . . . 4303.814 . . . 4314.162 4317.114 4319.650 4325.749 4326.457 4332.746 4336.832 4340.469 4345.535 4349.406 4351.269 4352.951 4359.528 4361.545 4363.195 4366.872 4368.442
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
S/N Ratio (9)
4121.463 4128.054 4128.748 4129.320 4130.894 4130.872 4131.63 4132.800 4143.759 4146.076 4153.298 4156.530 4168.972 4177.196 4178.958 4185.440 4189.788 4189.581 4189.593 4189.681 4189.70 4189.76 4201.172 ... 4211.099 4219.74 4236.91 4237.05 4241.246 4241.78 4243.969 4248.799 4253.499 4253.381 4267.001 4267.183 4267.261 4275.551 4276.749 4276.829 4284.904 4287.727 4287.394 4294.402 4294.782 4303.611 4303.823 4314.29 4317.139 4319.630 4325.761 4326.237 4332.653 4336.859 4340.450 4345.560 4349.426 4351.260 4352.778 4359.333 4361.468 4363.209 4366.895 4368.193
O II Si II Fe II O II Si II Si II N III O II He I O II O II O II He I [Fe II] [Fe II] O II O II O II S II S II Al III Al III [Ni II] ... [Fe II] Ne II N II N II Cl II N II [Fe II] [Ni II] S III Cl II C II C II C II O II O II [Fe II] S III O II [Fe II] S II O II O II O II Fe II O II O II O II [Ni II] S III O II HI O II O II O II [Fe II] [Fe II] S III [O III] O II OI
... ? ? ? ? ? ? ... ... ? ... ? ... ? ? ... ... ... ... ... ... ... ? ? ? ? ? ? ? ? ... ? ? ? ... ... ... ... ... ... ... ... ... ? ? ... ... ? ... ... ... ? ... ? ... ... ... ? ... ... ? ... ... ...
[4.0 [12.7 [12.6 [4.8 ... ... 14.7 [1.8 [0.4 [1.7 [1.5 [15.2 3.3 15.0 18.7 [0.1 ... ... ... ... ... ... 14.6 ... 13.5 0.5 ... ... 10.6 [1.3 12.4 16.9 ... ... ... ... ... 0.2 ... ... 1.9 ... ... [4.4 [0.3 ... ... [8.9 [1.7 1.4 [0.9 15.2 6.4 [1.9 1.3 [1.7 [1.3 0.6 11.9 13.4 5.3 [0.9 [1.6 ...
23 23 15 12 13 ... 38 14 17 19 14 16 26 16 5 17 14 ... ... ... ... ... 13 10 8 14 24 ... 17 17 17 13 13 ... 28 ... ... 13 20 ... 11 12 ... 10 13 16 ... 29 16 17 15 13 27 13 26 15 16 15 19 13 23 14 14 11
0.0057 0.0015 0.0017 0.0014 0.0012 ... 0.0040 0.0049 0.0552 0.0016 0.0068 0.0079 0.0092 0.0016 0.0006 0.0033 0.0039 ... ... ... ... ... 0.0017 0.0007 0.0005 0.0010 0.0015 ... 0.0028 0.0013 0.0074 0.0014 0.0029 ... 0.0449 ... ... 0.0025 0.0058 ... 0.0023 0.0147 ... 0.0008 0.0012 0.0031 ... 0.0031 0.0059 0.0035 0.0020 0.0058 0.0055 0.0021 10.6472 0.0085 0.0102 0.0021 0.0033 0.0103 0.0020 0.2148 0.0051 0.0095
0.0093 0.0025 0.0028 0.0023 0.0020 ... 0.0065 0.0080 0.0900 0.0026 0.0111 0.0128 0.0149 0.0026 0.0010 0.0053 0.0063 ... ... ... ... ... 0.0027 0.0011 0.0008 0.0016 0.0024 ... 0.0045 0.0021 0.0118 0.0022 0.0046 ... 0.0712 ... ... 0.0040 0.0092 ... 0.0036 0.0232 ... 0.0013 0.0019 0.0049 ... 0.0049 0.0093 0.0055 0.0031 0.0091 0.0086 0.0033 16.5986 0.0132 0.0159 0.0033 0.0051 0.0160 0.0031 0.3332 0.0079 0.0147
22.2 5.1 8.4 7.7 6.2 ... 8.7 23.0 214.4 5.6 31.5 35.4 26.4 5.6 4.7 9.0 12.0 ... ... ... ... ... 9.0 4.6 4.4 7.0 6.6 ... 14.7 6.8 37.6 7.1 11.7 ... 85.1 ... ... 10.6 18.1 ... 10.0 64.0 ... 5.5 6.9 15.2 ... 9.5 28.3 15.8 8.9 30.0 12.8 7.6 23204.7 27.0 29.2 6.1 10.6 18.6 6.3 964.7 22.8 48.9
233
Notes (10) same multiplet as 4130.884
same multiplet as 4128.054 large FWHM
possible blend with C III 4156.504 blend anomalously low FWHM average average
same multiplet as 4326.237 anomalously low FWHM
possible blend with O II 4253.74, 4253.98 from the same multiplet as 4241.246 average, blend
same multiplet as 4253.499, 4332.653, 4361.468 possible recombination line
large FWHM possible blend with [Fe II] 4319.62 line same multiplet as 4201.172 gh average average average average
TABLE 1ÈContinued Rest Wavelength (A ) (1)
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
S/N Ratio (9)
OI OI He I Ne II Ne II [Fe II] O II [Fe II] O II Ti II Ca I He I O II [Fe II] [Fe II] S II O II N II O II O II He I [Fe II] O II Mg II Mg II Mg II S II O II [Fe II] [Fe II] Si III Mg I O II O II [Ni III] N II O II [Fe III] N II O II N II Si II Si II Si II [C I] N II [Ni II] N II N III N III O II N III O II N II O II
... ... ... ... ... ... ... ... ... ? ? ... ... ... ... ? ... ... ? ... ... ... ? ? ? ? ? ... ? ? ? ... ... ... ... ... ... ... ... ... ? ... ... ... ... ... ? ... ... ? ... ... ... ... ...
... ... 0.0 [1.3 [1.4 13.8 [0.3 13.1 [0.3 4.1 7.5 0.2 ... ... 13.6 0.2 ... ... [0.3 3.7 0.8 13.4 [2.6 ... ... ... 1.6 1.0 13.2 8.5 [3.5 6.9 [0.2 [0.1 13.6 [0.4 [3.2 ... ... 0.3 [1.3 ... ... ... ... ... 16.1 ... ... [3.8 [0.9 [4.6 [0.3 0.0 [0.1
... ... 17 16 13 11 13 15 14 23 22 16 13 ... 17 12 15 ... 11 17 19 14 28 14 14 14 11 16 11 16 10 19 13 14 18 15 20 14 ... 19 9 21 ... ... ... ... 9 14 ... 13 14 13 14 15 14
... ... 0.1026 0.0022 0.0017 0.0076 0.0056 0.0092 0.0046 0.0012 0.0015 0.0127 0.0059 ... 0.0032 0.0010 0.0024 ... 0.0010 0.0014 0.8535 0.0023 0.0009 0.0005 ... ... 0.0006 0.0018 0.0009 0.0009 0.0008 0.0014 0.0040 0.0032 0.0019 0.0025 0.0015 0.0094 ... 0.0023 0.0015 0.0030 ... ... ... ... 0.0008 0.0076 ... 0.0008 0.0083 0.0013 0.0166 0.0037 0.0266
... ... 0.1581 0.0034 0.0026 0.0116 0.0086 0.0141 0.0070 0.0018 0.0023 0.0194 0.0090 ... 0.0049 0.0015 0.0036 ... 0.0015 0.0021 1.2895 0.0035 0.0014 0.0008 ... ... 0.0009 0.0027 0.0014 0.0013 0.0012 0.0021 0.0059 0.0047 0.0028 0.0037 0.0022 0.0138 ... 0.0034 0.0022 0.0044 ... ... ... ... 0.0012 0.0111 ... 0.0012 0.0121 0.0019 0.0241 0.0054 0.0386
... ... 451.1 11.7 11.6 63.2 42.5 64.6 33.5 6.0 5.8 30.6 33.5 ... 15.7 6.6 14.0 ... 7.4 8.3 4822.3 16.1 4.0 4.5 ... ... 5.9 12.1 7.8 4.0 6.8 7.4 27.2 21.3 10.5 15.9 7.5 57.1 ... 9.9 8.2 7.3 ... ... ... ... 4.7 37.2 ... 5.2 53.9 9.0 109.0 22.5 171.5
O II [Fe III] O II [Fe III] O II O II O II O II
... ... ... ... ... ... ? ...
[0.7 6.4 [0.7 2.4 [0.8 [0.7 [2.3 ...
13 15 14 16 13 16 15 23
0.0081 0.1246 0.0091 0.0050 0.0014 0.0060 0.0011 0.0022
0.0117 0.1804 0.0132 0.0072 0.0020 0.0087 0.0016 0.0032
54.9 844.6 67.4 35.9 11.5 43.5 5.9 7.5
4634.072 4638.841 4640.569 4641.806 4643.086 4649.133
... ... ... ... ... ...
4368.242 4368.258 4387.929 4391.995 4409.30 4413.781 4414.899 4416.266 4416.975 4417.719 4425.437 4437.554 4452.378 4452.098 4457.945 4463.581 4465.407 4465.529 4467.924 4469.378 4471.486 4474.904 4477.906 4481.126 4481.150 4481.325 4483.427 4491.222 4492.634 4514.900 4567.840 4571.096 4590.974 4596.177 4596.84 4601.478 4602.129 4607.03 4607.153 4609.436 4613.868 4621.418 4621.696 4621.722 4621.570 4621.393 4628.046 4630.539 4630.61 4634.13 4638.856 4640.64 4641.810 4643.086 4649.135
4650.827 4658.149 4661.621 4667.047 4673.720 4676.224 4696.317 4699.167
... ... ... ... ... ... ... ...
4650.838 4658.05 4661.632 4667.01 4673.733 4676.235 4696.353 4699.011
4387.929 4391.976 4409.279 4413.984 4414.895 4416.459 4416.970 4417.779 4425.548 4437.556 4452.304
... ... ... ... ... ... ... ... ... ... ...
4458.148 . . . 4463.584 . . . 4465.424 . . . 4467.920 4469.433 4471.491 4475.104 4477.868 4481.281
... ... ... ... ... ...
4483.450 4491.238 4492.831 4515.027 4567.787 4571.200 4590.972 4596.175 4597.048 4601.472 4602.080 4607.114
... ... ... ... ... ... ... ... ... ... ... ...
4609.440 . . . 4613.848 . . . 4621.355 . . .
4628.294 . . . 4630.536 . . .
234
Notes (10)
same multiplet as 4409.300 same multiplet as 4391.995 same multiplet as 4416.975 same multiplet as 4414.899 strongest line in multiplet average
strongest line in multiplet
same multiplet as 4465.407 same multiplet as 4465.407 gh
from multiplet with 11 lines, strongest 4416.266
the strongest semiforbidden optical line
average ; anomalously low FWHM average observed line might include Si II triplet
main contributor anomalously low FWHM
from multiplet with strongest line at 4640.64 from multiplet with strongest line at 4640.64
this is the strongest line of the multiplet ; all lines are observed
cosmic ray hit on wing weakest line from multiplet 4649.135
TABLE 1ÈContinued Rest Wavelength (A ) (1)
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
4699.218 4701.53 4705.346 4711.37 4713.139 4728.068 4733.91 4740.17 4752.691 4754.69 ... ... 4769.43 4774.718 4777.68 4779.722 4788.138 4789.45 4792.007 4793.648 ... 4803.287
O II [Fe III] O II [Ar IV] He I [Fe II] [Fe III] [Ar IV] O II [Fe III] ... ... [Fe III] [Fe II] [Fe III] N II N II [F II] S II N II ... N II
... ... ... ... ... ? ... ... ? ... ? ? ... ... ... ... ... ? ? ? ? ...
... 4.6 0.8 [2.5 1.5 14.3 1.7 1.1 16.5 5.4 ... ... 4.2 14.1 3.4 [0.4 0.0 10.1 4.2 12.0 ... [1.2
... 14 13 13 19 14 15 13 22 15 21 21 15 13 14 14 17 16 23 25 15 12
... 0.0390 0.0030 0.0109 0.1340 0.0011 0.0166 0.0125 0.0016 0.0232 0.0021 0.0010 0.0136 0.0012 0.0080 0.0014 0.0020 0.0009 0.0010 0.0019 0.0017 0.0023
... 0.0559 0.0043 0.0156 0.1916 0.0016 0.0236 0.0178 0.0023 0.0329 0.0030 0.0014 0.0192 0.0017 0.0113 0.0020 0.0028 0.0013 0.0014 0.0027 0.0024 0.0032
... 188.8 14.2 38.9 354.3 7.0 100.8 92.5 8.1 161.6 11.1 5.6 88.4 8.3 53.4 9.6 12.1 5.3 4.3 7.7 8.6 13.2
4814.741 . . .
4814.534
[Fe II]
...
12.9
18
0.0086
0.0120
28.5
4815.563 . . .
4815.617 4815.552 4861.33 4867.12 4867.17 4867.066 4881.00 4889.617
N II S II Hb N III N III C II [Fe III] [Fe II]
... ... ... ... ... ... ... ...
... ... 0.2 ... ... ... 4.1 13.3
17 ... 26 31 ... ... 15 13
0.0030 0.0024 17.4964 0.0047 ... ... 0.0587 0.0049
0.0042 ... 24.2016 0.0065 ... ... 0.0809 0.0067
10.3 ... 62680.0 14.1 ... ... 434.4 29.4
4895.117 4905.339 4921.931 4924.50 4924.529 4930.54 4931.227 4947.373 4950.744 4958.911 4973.388 4987.376 4994.360 4994.370 5006.843 5011.26 5015.678 5018.440 5035.399 5041.024 5047.738 5055.984 5056.317 5084.77 ... 5111.627 5121.828 ...
N II [Fe II] He I [Fe III] O II [Fe III] [O III] [Fe II] [Fe II] [O III] [Fe II] N II N II N II [O III] [Fe III] He I Fe II [Fe II] Si II He I Si II Si II [Fe III] ... [Fe II] C II ...
? ? ... ... ... ... ... ... ... ... ? ... ? ? ... ... ... ? ... ... ... ... ... ... ? ... ... ?
[11.6 13.2 0.4 ... ... 5.9 1.5 11.2 17.8 1.2 16.0 [6.0 ... ... 0.6 4.5 0.1 [6.7 19.1 1.2 0.9 ... ... 2.8 ... 13.3 1.7 ...
19 12 17 16 ... 19 14 16 34 13 20 16 12 ... 13 13 16 14 23 18 18 22 ... 15 11 12 17 19
0.0029 0.0017 0.2596 0.0046 ... 0.0055 0.0094 0.0036 0.0026 24.7511 0.0027 0.0137 0.0024 ... 74.1110 0.0164 0.4993 0.0063 0.0124 0.0198 0.0394 0.0415 ... 0.0032 0.0005 0.0029 0.0018 0.0013
0.0040 0.0023 0.3544 0.0063 ... 0.0075 0.0128 0.0049 0.0035 33.5152 0.0036 0.0185 0.0032 ... 99.3964 0.0220 0.6683 0.0084 0.0165 0.0264 0.0524 0.0551 ... 0.0042 0.0007 0.0038 0.0024 0.0017
10.7 6.7 616.2 12.8 ... 15.5 32.5 22.1 8.4 79477.3 3.5 31.9 4.2 ... 205787.5 39.8 830.5 7.5 20.1 27.9 57.1 75.8 ... 24.7 7.0 27.5 11.3 6.0
4701.603 4705.358 4711.331 4713.164 4728.294 4733.937 4740.188 4752.953 4754.776 4756.469 4762.400 4769.496 4774.942 4777.734 4779.715 4788.139 4789.611 4792.073 4793.839 4802.489 4803.267
4861.334 . . . 4867.063 . . .
4881.066 . . . 4889.833 . . . 4894.927 4905.555 4921.937 4924.524
... ... ... ...
4930.638 4931.252 4947.557 4951.038 4958.930 4973.653 4987.276 4994.377
... ... ... ... ... ... ... ...
5006.854 5011.335 5015.679 5018.328 5035.720 5041.044 5047.753 5056.027
... ... ... ... ... ... ... ...
5084.817 5089.326 5111.854 5121.856 5131.924
... ... ... ... ...
235
S/N Ratio (9)
Notes (10)
average average possible line from multiplet with 4889.617
from multiplet with strongest line at 4814.534
several lines of this multiplet are observed the same multiplet as 4774, 4779, 4781, 4788, 4793, 4810 average ; from multiplet with strongest line at 4814.534 average ... large FWHM strongest line in multiplet (blend of two)
but possibly [Fe II] 4889.70 from a di†erent multiplet recombination line, singlet average
large FWHM measured from short exposures same multiplet as 4994.37
short exposures short exposures the same multiplet as 5169.033
B B B
TABLE 1ÈContinued Rest Wavelength (A ) (1)
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
S/N Ratio (9)
5146.890 5146.917 5158.951 5158.956 5169.274
... ... ... ... ...
5146.749 5146.749 5158.777 5158.777 5169.033
Co I Co I [Fe II] [Fe II] Fe II
? ? ... ... ...
8.2 9.8 10.1 10.4 14.0
17 18 21 21 13
0.0049 0.0064 0.0140 0.0113 0.0014
0.0064 0.0083 0.0182 0.0147 0.0018
8.9 25.5 56.7 56.6 10.8
5169.288 5191.700 5198.159 5198.169 5200.516 5200.526 5219.353 5219.359 5220.275 5220.297 5261.854 5261.856 5269.228 5270.530
... ... ... ... ... ... ... ... ... ... ... ... ... ...
5169.033 5191.816 5197.902 5197.902 5200.257 5200.257 5219.307 5219.307 5220.059 5220.059 5261.621 5261.621 5268.874 5270.40
Fe II [Ar III] [N I] [N I] [N I] [N I] S III S III [Fe II] [Fe II] [Fe II] [Fe II] [Fe II] [Fe III]
... ... ... ... ... ... ? ... ? ... ... ... ? ...
14.8 [6.7 14.8 15.4 14.9 15.5 2.6 3.0 12.4 13.7 13.3 13.4 20.2 7.4
10 14 11 11 12 11 18 14 19 20 15 13 20 15
0.0011 0.0127 0.0203 0.0253 0.0119 0.0155 0.0007 0.0009 0.0011 0.0012 0.0107 0.0088 0.0007 0.0638
0.0014 0.0164 0.0262 0.0327 0.0154 0.0200 0.0009 0.0012 0.0014 0.0015 0.0136 0.0112 0.0009 0.0812
8.8 63.7 291.8 287.0 168.5 175.3 3.9 9.8 5.5 9.9 38.2 71.3 4.1 335.1
5270.543 . . .
5270.40
[Fe III]
...
8.1
15
0.0649
0.0826
252.8
5273.577 . . . 5273.598 . . . 5275.306 . . .
5273.346 5273.346 5275.123 5275.166 5275.123 5275.166 5296.829 5299.044 5299.088 5299.044 5299.088 5333.646 5333.646 ... ... 5363.34 5363.34 5376.452 5376.452 5411.98 5411.98 5432.797 5432.797 5433.129 5433.129 5453.855 5453.855 5495.655 5495.655
[Fe II] [Fe II] OI OI OI Oi [Fe II] OI OI OI OI [Fe II] [Fe II] ... ... [Ni IV] [Ni IV] [Fe II] [Fe II] [Fe III] [Fe III] S II S II [Fe II] [Fe II] S II S II N II N II
... ... ... ... ? ? ? ... ... ... ... ... ... ? ? ? ? ... ... ... ... ? ? ... ... ... ... ... ...
13.2 14.3 ... ... ... ... 11.8 ... ... ... ... 12.2 12.4 ... ... 17.9 17.9 8.6 10.8 8.3 9.4 [1.9 1.5 13.0 13.3 [0.7 [0.1 [1.6 [0.8
16 16 19 ... 20 ... 15 16 ... 18 ... 17 18 14 15 15 11 40 23 18 13 14 15 15 20 13 11 24 15
0.0057 0.0053 0.0018 ... 0.0015 ... 0.0006 0.0046 ... 0.0055 ... 0.0028 0.0025 0.0028 0.0026 0.0014 0.0008 0.0037 0.0023 0.0068 0.0062 0.0009 0.0012 0.0018 0.0017 0.0013 0.0012 0.0019 0.0013
0.0073 0.0067 0.0023 ... 0.0019 ... 0.0008 0.0058 ... 0.0070 ... 0.0035 0.0031 0.0035 0.0033 0.0017 0.0010 0.0046 0.0029 0.0084 0.0077 0.0011 0.0015 0.0022 0.0021 0.0016 0.0015 0.0023 0.0016
35.3 26.4 8.5 ... 6.3 ... 4.7 36.1 ... 37.2 ... 20.9 16.5 26.5 26.2 10.4 8.4 8.4 12.1 30.3 42.9 6.9 6.2 7.6 10.1 9.1 11.4 9.2 14.9
5512.978 . . .
5506.778 5507.00 5512.772
Fe I SI OI
? ? ...
... ... ...
10 ... 16
0.0005 ... 0.0035
0.0006 ... 0.0043
5.3 ... 40.2
5512.994 . . .
5512.820 5512.980 5512.772
OI Ca I OI
... ... ...
... ... ...
... ... 16
... ... 0.0038
... ... 0.0046
... ... 14.3
5512.820 5512.980
OI Ca I
... ...
... ...
... ...
... ...
... ...
... ...
5275.339 . . . 5297.037 . . . 5299.245 . . . 5299.252 . . . 5333.862 5333.866 5342.392 5342.395 5363.659 5363.665 5376.606 5376.646 5412.130 5412.149 5432.762 5432.824 5433.365 5433.369 5453.843 5453.853 5495.626 5495.640
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
5506.859 . . .
236
Notes (10) R B B R R possibly pumped ; strongest line in multiplet ; other possible 5018.440 B R average R B R B R B R average B B possible contribution from Ca I 5261.704 3DÈ3 P R possible contribution from Ca I 5261.704 3DÈ3 P R R average ; possible contribution from Ca I 5270.270 3DÈ 3P B average ; possible contribution from Ca I 5270.270 3DÈ 3P R average B R average R B cannot see on two-dimensional image B R R R B B B R B R B R B same multiplet as 5159, 5262 R average ; fuzzy on two-dimensional image B R B R same multiplet as 5453.855 R average ; same multiplet as 5273 B R same multiplet as 5432.797 B B, gh R several weak lines of this multiplet possibly observed R R R average ; ID lambda is average of sextet with all lines within 0.2A R R B wavelength is approximate, averaged for several lines in the multiplet B B
TABLE 1ÈContinued Rest Wavelength (A ) (1)
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
5517.675 . . .
5517.66 5517.66 5517.66 5517.66 5518.102 5518.576 5518.102 5518.576 ... 5535.347 5535.353 5535.383
[Cl III] [Ni IV] [Cl III] [Ni IV] NI NI NI NI ... N II C II N II
... ... ... ... ? ? ? ? ... ? ? ?
... ... ... ... ... ... ... ... ... ... ... ...
14 ... 14 ... 14 ... 14 14 21 24 ... ...
0.0810 ... 0.0832 ... 0.0047 ... 0.0050 ... 0.0027 0.0012 ... ...
0.0984 ... 0.1011 ... 0.0057 ... 0.0061 ... 0.0033 0.0015 ... ...
456.3 ... 258.4 ... 8.9 ... 26.9 ... 16.7 5.3 ... ...
5537.829 . . . 5537.830 . . . 5551.843 . . .
5537.7 5537.7 5551.922
[Cl III] [Cl III] N II
... ... ...
... ... [4.3
17 14 20
0.1257 0.1220 0.0014
0.1522 0.1477 0.0017
701.9 856.1 8.6
5551.901 . . . 5555.209 . . .
5551.922 5555.004 5555.053 5555.004 5555.053 5577.339 5577.339 5606.151 5666.63 5666.63 5676.02 5676.02 5679.56 5679.56 5686.21 5686.21 5710.77 5710.77 5739.73 5739.73 5746.966 5754.59 5754.59 ... 5867.6 5867.8 5867.90 5875.64 5875.64 ... 5889.28 5889.78 5891.60 ... 5927.81 5927.81 5931.78 5931.78 5941.65 5952.39 5957.56 5958.386 5958.584 5958.640 5978.93 6000.2
N II OI OI OI OI [O I] [O I] S II N II N II N II N II N II N II N II N II N II N II Si III Si III [Fe II] [N II] [N II] ... Al II Al II Al II He I He I ... C II C II C II ... N II N II N II N II N II N II Si II OI OI OI Si II [Ni III]
... ... ... ... ... ... ... ? ? ... ... ... ... ... ? ... ... ... ... ... ? ... ... ? ? ? ? ... ... ? ? ? ... ? ... ? ... ... ... ? ... ... ... ... ... ...
[1.1 ... ... ... ... 9.9 10.7 [0.3 [1.3 [0.6 [1.6 [1.1 [1.0 [0.5 [1.7 [0.6 [1.8 [0.7 [2.0 [1.6 11.7 3.8 3.9 ... ... ... ... [0.6 [0.5 ... ... ... [1.9 ... 0.3 1.0 [1.6 1.2 [1.4 0.4 5.7 ... ... ... 5.9 1.1
19 19 ... 17 ... 22 20 11 15 15 25 15 18 15 17 13 18 13 13 14 17 18 19 11 32 ... ... 22 19 23 26 ... 17 17 13 14 14 16 12 25 23 19 ... ... 24 14
0.0014 0.0055 ... 0.0050 ... 0.0028 0.0028 0.0005 0.0045 0.0063 0.0038 0.0024 0.0088 0.0088 0.0018 0.0016 0.0020 0.0018 0.0047 0.0049 0.0007 0.1727 0.1693 0.0011 0.0017 ... ... 3.0441 3.1916 0.0011 0.0026 ... 0.0017 0.0018 0.0016 0.0017 0.0031 0.0029 0.0026 0.0013 0.0109 0.0067 ... ... 0.0242 0.0027
0.0017 0.0066 ... 0.0060 ... 0.0034 0.0034 0.0006 0.0053 0.0075 0.0045 0.0028 0.0104 0.0104 0.0021 0.0019 0.0023 0.0021 0.0055 0.0057 0.0008 0.2014 0.1974 0.0013 0.0019 ... ... 3.4745 3.6428 0.0013 0.0030 ... 0.0019 0.0020 0.0018 0.0019 0.0035 0.0033 0.0029 0.0015 0.0123 0.0075 ... ... 0.0271 0.0030
8.1 32.2 ... 24.8 ... 16.4 15.6 6.0 6.3 55.3 9.5 21.9 30.2 80.1 7.1 15.0 12.0 14.3 44.9 42.3 5.5 974.0 1162.2 10.6 6.9 ... ... 17593.2 18100.6 6.2 13.3 ... 11.6 11.6 16.9 7.0 31.3 8.8 27.4 6.1 51.0 37.1 ... ... 125.9 29.3
5517.680 . . . 5518.354 . . . 5518.342 . . . 5527.508 . . . 5535.357 . . .
5555.224 . . . 5577.523 5577.537 5606.145 5666.606 5666.619 5675.990 5676.000 5679.542 5679.551 5686.178 5686.198 5710.736 5710.757 5739.691 5739.699 5747.190 5754.663 5754.664 5821.106 5867.806
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
5875.628 5875.630 5887.744 5889.506
... ... ... ...
5891.564 5906.082 5927.817 5927.830 5931.748 5931.803 5941.622 5952.398 5957.673 5958.795
... ... ... ... ... ... ... ... ... ...
5979.047 . . . 6000.222 . . .
237
S/N Ratio (9)
Notes (10) R average, blended with weak line R B blended with weak line B B average, blend with strong line B average R R average, blend with strong line R not seen in blue spectrum R R R several weak lines of this multiplet are possibly observed B R average R several weak lines of this multiplet are possibly observed B B B R average R B R B strongest line in multiplet B R B R B R B R B R R B B R average B R average B probably not real ; large FWHM B B B R average B B includes night sky B includes night sky B B night sky ? R B R blended with sky B
average, blend average
average, blend
TABLE 1ÈContinued Rest Wavelength (A ) (1)
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
S/N Ratio (9)
6046.675 . . .
OI OI ... Ne I ... [O I] [S III] Si II [O I] [Ni II] Si II [Ni III] Ne I Ne I [Fe II] ... [Ni III] [N II] ... Ha C II [N II] [Ni II] He I [Ni III] [S II] [S II] ... C III [Cr IV] C II C II
... ... ? ? ? ... ... ... ... ... ... ... ... ? ? ... ? ... ? ... ... ... ... ... ? ... ... ? ? ? ? ?
... ... ... 0.2 ... 11.0 0.7 1.8 11.1 ... 1.2 ... ... 1.3 5.7 ... ... 3.3 ... [0.7 [1.2 2.1 14.2 [0.2 ... 4.4 4.8 ... [8.7 ... [0.4 [0.3
18 ... 17 16 26 17 14 20 16 10 19 13 ... 13 19 14 13 17 57 27 14 17 12 17 15 20 19 14 15 12 14 9
0.0167 ... 0.0005 0.0006 0.0025 0.1853 0.4755 0.0420 0.0621 0.0020 0.0212 0.0014 ... 0.0018 0.0005 0.0060 0.0038 3.6260 0.0849 73.7560 0.0582 11.3392 0.0035 1.0000 0.0014 0.5331 1.0041 0.0009 0.0009 0.0010 0.0010 0.0005
0.0185 ... 0.0005 0.0007 0.0027 0.1973 0.5050 0.0444 0.0654 0.0021 0.0223 0.0015 ... 0.0019 0.0005 0.0062 0.0039 3.7086 0.0868 75.2211 0.0592 11.5179 0.0035 ... 0.0014 0.5307 0.9969 0.0009 0.0009 0.0010 0.0010 0.0005
92.6 ... 3.9 5.2 12.5 1028.1 2908.8 227.4 469.5 20.9 153.1 13.5 ... 17.7 4.6 39.2 11.4 9685.7 109.8 6091.4 273.2 2297.1 29.2 7808.0 12.8 3263.8 6625.0 7.7 5.0 9.0 10.7 6.7
6139.321 6143.066 6151.329 6300.536 6312.075 6347.147 6364.011 6365.433 6371.396 6401.230
... ... ... ... ... ... ... ... ... ...
6402.273 6440.522 6461.836 6533.560 6548.123 6549.770 6562.804 6578.024 6583.496 6667.116 6678.146 6681.921 6716.539 6730.924 6739.860 6744.193 6747.632 6779.932 6780.584
... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
6046.438 6046.495 ... 6143.063 ... 6300.304 6312.06 6347.11 6363.776 6365.10 6371.37 6401.5 6401.076 6402.246 6440.400 ... 6533.8 6548.05 ... 6562.82 6578.05 6583.45 6666.80 6678.152 6682.2 6716.44 6730.816 ... 6744.39 6747.5 6779.94 6780.59
6813.919 6826.882 6827.961 6829.845
... ... ... ...
6813.57 ... 6828.12 6829.80
[Ni II] ... CI Si II
? ? ? ?
15.4 ... [7.0 ...
11 18 16 13
0.0015 0.0017 0.0016 0.0007
0.0015 0.0017 0.0016 0.0007
15.4 10.5 10.6 4.4
6855.978 6861.626 6890.552 6906.403 6910.534 6915.437 6933.958
... ... ... ... ... ... ...
6829.83 ... ... ... 6906.436 6910.562 6915.2 6933.910 6933.661 6946.4 6989.470 7002.173
Si II ... ... ... O II O II [Cr IV] He I [Fe II] [Ni III] He I OI
... ? ? ? ... ? ? ... ... ? ... ...
... ... ... ... [1.4 [1.2 ... ... ... ... 2.6 ...
... 14 10 16 14 15 12 15 ... 12 20 24
... 0.0015 0.0006 0.0019 0.0013 0.0008 0.0009 0.0027 ... 0.0008 0.0032 0.0194
... 0.0015 0.0006 0.0018 0.0013 0.0008 0.0009 0.0026 ... 0.0008 0.0030 0.0184
... 13.3 7.1 14.7 12.0 7.0 8.3 14.3 ... 4.0 17.7 91.8
7002.196 7002.230 7002.250 7032.413 7062.280 7065.179 7065.217 ... 7135.790 7155.16
OI OI OI Ne I He I He I He I ... [Ar III] [Fe II]
... ... ... ? ... ... ... ? ... ...
... ... ... [0.2 2.0 ... ... ... [0.7 8.6
... ... ... 10 17 18 ... 20 13 19
... ... ... 0.0005 0.0050 2.1055 ... 0.0008 4.6787 0.0162
... ... ... 0.0005 0.0047 1.9813 ... 0.0008 4.3507 0.0150
... ... ... 6.4 25.2 10405.8 ... 5.2 38291.5 76.4
6946.131 . . . 6989.531 . . . 7002.436 . . .
7032.408 . . . 7062.328 . . . 7065.203 . . . 7080.675 . . . 7135.773 . . . 7155.366 . . .
Notes (10) average
blend average average average average average average ; blend
identiÐed by Esteban 1998 average average ; same multiplet as 6946.4 average average ; on wing of [N II] ; large FWHM average average average average average average average average average average ; strongest in multiplet average ; strongest in multiplet same multiplet as 6780.59 same multiplet as 6779.94, but strongest line (6783.91) is missing ; anomalously low FWHM average average ; measured after subtracting sky average ; measured after subtracting sky average ; after subtracting sky ; doublet of Si II average after subtracting sky average average average average
same multiplet as 4596.84, 6533.8 average average, blend or cut by atmospheric absorption
average average
average
FAINT EMISSION LINES IN THE ORION NEBULA
239
TABLE 1ÈContinued Rest Wavelength (A ) (1)
ID Wavelength (A ) (2)
ID (3)
ID ? (4)
dV (km s~1) (5)
FWHM (km s~1) (6)
I/6678 obs (7)
I/6678 cor (8)
S/N Ratio (9)
7160.559 . . . 7172.218 . . .
7160.600 7172.00
He I [Fe II]
... ?
[1.7 9.1
16 13
0.0073 0.0038
0.0068 0.0035
41.9 6.5
7231.286 . . . 7236.424 . . . 7237.102 . . .
7231.33 7236.42 7237.17
C II C II C II
? ? ?
[1.8 0.2 [2.8
17 14 22
0.0162 0.0506 0.0099
0.0148 0.0463 0.0091
103.5 359.4 22.3
7254.728 . . .
7254.448
OI
...
...
22
0.0256
0.0234
91.2
7281.355 7298.019 7319.106 7320.184 7329.728 7330.822 7378.124 7388.397 7411.994 7424.002 7452.782 7468.689
7254.531 7281.351 7298.050 7318.92 7319.99 7329.67 7330.73 7377.83 7388.18 7411.61 7423.641 7452.54 7468.312
OI He I He I [O II] [O II] [O II] [O II] [Ni II] [Fe II] [Ni II] NI [Fe II] NI
... ... ... ... ... ... ... ... ... ... ... ... ...
... 0.2 [1.3 7.6 8.0 2.4 3.8 12.0 8.8 15.5 14.6 9.7 15.1
... 20 18 18 17 16 18 20 20 13 11 19 13
... 0.1684 0.0099 0.4923 1.4107 0.7478 0.7832 0.0247 0.0027 0.0070 0.0031 0.0051 0.0120
... 0.1529 0.0090 0.4450 1.2749 0.6746 0.7064 0.0221 0.0024 0.0062 0.0028 0.0045 0.0106
... 747.4 69.0 2676.0 7986.8 4987.6 4958.1 155.3 11.2 72.2 42.3 39.2 105.0
... ... ... ... ... ... ... ... ... ... ... ...
The results are given in Table 1. The columns are as follows. 1. Observed wavelength in the Orion rest frame, in A ; 2. laboratory wavelength, in A ; 3. identiÐcation ; 4. ““ ? ÏÏ for lines with no identiÐcation, or with small signal-to-noise ratio, or which are otherwise uncertain ; 5. velocity relative to Orion rest frameÈvelocities are not determined in the case of blends and for several lines with inaccurate laboratory wavelength ; 6. FWHM ; 7. ratio of the observed intensity to the intensity of the He I j6678 line ; 8. reddening-corrected line ratios relative to the intensity of He I j6678 ; 9. signal-to-noise ratio ; and 10. notesÈlines within the overlapped part are marked by B (blue spectrum) or R (red spectrum). We initially measured everything that appeared to be an emission line in each order of the one-dimensional co-added red and blue spectra. We then went back and inspected the position of each line in the original two-dimensional images, which allowed us to exclude from the line list any night sky lines, remnants of cosmic-ray hits, and artifacts caused by scattered light, ghosts, or internal reÑections in the spectrograph. The superimposed sky lines are obviously narrower than the Orion lines and have a more uniform intensity along the slit, thus we could check against our o†set sky spectrum as needed. On the two-dimensional image, cosmic-ray hits are easily recognized as round spots,
Notes (10) average Ñux uncertain ; cut by atmospheric absorption cut by atmospheric absorption cut by atmospheric absorption lambda uncertain, blend with strong line ; same multiplet as 7236.42 ; 7231.33 average, blend or cut by atmospheric absorption average cut by atmospheric absorption
while most ghosts and reÑections are either fuzzy or do not fall exactly on the echelle orders. However, the blue spectrum also included two detectable pairs of Rowland ghosts spaced symmetrically around each of the strong emission lines and looking exactly like real emission lines. These are a feature of the 79 line mm~1 Bausch & Lomb echelle grating used in this spectrograph (and in many other spectrographs). One pair of these ghosts has intensities 0.00135(5007/j)4 times that of the real emission line and is separated from the real line by approximately ^2.000(j/5007)2 A , while the second pair has an intensity 0.00045(5007/j)4 times the real line and separation ^4.000(j/5007)2A . These artifacts were removed from the blue spectrum by Ðrst convolving each order with a kernel that would reproduce the ghost spectrum and then subtracting the ghost spectrum. A few real emission lines were blended with these ghosts, and their intensities were measured after subtracting o† the ghosts. Those lines are identiÐed in Table 1 by the notation ““ gh.ÏÏ A few of the strongest emission lines were saturated on the 1000 s exposures, and therefore were measured from shorter exposures (100 s in the blue, 10 s in the red). Their measured wavelengths and Ñuxes were tied into the measurements from the 1000 s exposures by using lines of intermediate strength measured in both the long and short exposure data. The lines a†ected are [O III] jj4959, 5007 in the blue, and Ha, and [N II] jj6548, 6584 in the red. The S/N of each detected emission line was estimated by Ðrst measuring the rms pixel-to-pixel scatter in the continuum (p ), in units of Ñux per pixel, and then comparing the con line intensity to the noise level, taking into account both the
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BALDWIN ET AL.
measurement error based on the number of pixels across the line and the error in determining the continuum level. For the S/N in the emission-line Ñux, we used the relation S/N \ I /p line line \ I /p J[n (1 ] n /n )] . (1) line con line line con Here n is the number of pixels included in the line line proÐle (estimated as 2 ] FWHM in units of pixels), and n \ 10 is typical of the number of pixels averaged con together to determine the continuum level. This expression should be approximately correct for weak lines where the continuum noise dominates, and we checked for a few test cases so that it gives reasonable agreement with the S/N calculated from the scatter about Gaussian Ðts to the emission lines. Visual inspection of plots of the actual spectra showed that lines having S/N \ 3.5 are at the noise limit, while lines having S/N [ 7 are clearly detected. We therefore rejected all lines with S/N \ 3.5 and kept lines with 3.5 \ S/N \ 7 in our line list but with a ““ ? ÏÏ notation in column (4) of Table 1. The calculated S/N is given in column (8). Note that for strong lines the S/N is grossly overestimated, but this does not matter because we are interested only in having an objective criterion for including or excluding lines at the faint limit. In cases where an emission line was detected in 2 orders, the entry in Table 1 was computed by taking the average value of each measured quantity (including S/N) if either of the two individual signal-to-noise measurements were within a factor of 2 of each other or both had S/N [ 24. Such cases are marked ““ average ÏÏ in the Notes to Table 1(col. [10]). If these criteria were not met, we just used the measurements having higher S/N. The reason for taking an average rather than weighting by S/N is that the di†erences between measurements are more likely to be dominated by systematic errors (particularly in the Ñux calibration) than by noise. 3.
Vol. 129
The blue spectrum has approximately the same internal accuracy as the red spectrum. Comparing lines measured in 2 orders in the blue spectrum (35 cases), we Ðnd that the 1 p velocity scatter is ^1.5 km s~1 and the mean Ñux ratio is 0.97 with a 1 p scatter of ^0.24. We can also make an external check. Table 2 compares our new measurements to previous results for some of these lines measured at this same location using both the same echelle spectrograph (B96) and at nearby positions using the low-resolution Ritchey-Chretien (RC) spectrograph (BFM) on the Blanco Telescope. The agreement with the previous echelle data is in fact very good ; comparing the line strengths normalized to He I j6678, the mean ratio between the new and previous results is 1.01, with a 1 p scatter of ^0.05. This shows the measurements at the same location using the same instrumentation and calibration method are reproducible at a 5% level of accuracy. However, comparing the low-resolution spectra taken at nearly the same position (BFM, position 2), the mean is 1.06 with a 1 p scatter of ^0.14 (leaving out of this tabulated comparison [O I] jj5577, 6300, and 6363, which are blended with night sky lines in the low-resolution data, and
ANALYSIS OF ERRORS
Here we assess the uncertainties in the present measurements Ðrst by studying the internal consistency of the data set and then by comparison to previous results. Particularly in the case of the blue spectrum, the echelle Ñux calibration is liable to have been a†ected by scattered light (grating ghosts plus internal reÑections in the spectrograph optics) and by crowding of the orders in the UV. In the wavelength range where the blue and red spectra overlapped, the mean blue/red Ñux ratio is 1.11, but this is strongly inÑuenced by noise in the measurements of the many weak lines included in this average. Using only the Ðve strongest lines, the mean ratio is 1.00, while for He I j5876 (by far the strongest line in the overlap region) the ratio is 0.95. We concluded that the Ñux calibrations of the blue and red spectra are in good agreement. In the red spectrum, there are 37 cases where the same nebular emission line is recorded in more than 1 order. The 1 p velocity scatter between these duplicate measurements is ^1.5 km s~1, and the mean ratio of strengths is 1.0, with a 1 p scatter of ^0.2. This method depends on lines appearing in 2 orders and so tends to compare measurements taken far o† the echelle blaze, where the Ñux calibration is worst. It therefore overestimates the typical error. By contrast, intensities of intermediate-strength lines measured from exposures of di†erent lengths agree to about 3%, but in this latter comparison, Ñux calibration errors cancel out.
FIG. 3.ÈSpectrum near 5577 A showing the separation between the strong telluric emission and the nebular [O I] j5577 emission line (Orion rest frame).
FIG. 4.ÈLog of the O0 to O2` ionization fractions vs. depth in cloud
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FAINT EMISSION LINES IN THE ORION NEBULA
241
TABLE 2 RATIOS OF I(LINE)/I(6678) PREVIOUS OBSERVATIONS/THIS WORK RC (BFM) LINE
ECHELLE (B96)
Position 1
Position 2
Position 3
[O II] 3727 . . . . . . . . . . . . . . . . . . . . . . . . . [Ne III] 3869 . . . . . . . . . . . . . . . . . . . . . . . Blend 4340 ] 4363 . . . . . . . . . . . . . . . H I 4340 . . . . . . . . . . . . . . . . . . . . . . . . . . . . [O III] 4363 . . . . . . . . . . . . . . . . . . . . . . . . He I 4471 . . . . . . . . . . . . . . . . . . . . . . . . . . . H I 4861 . . . . . . . . . . . . . . . . . . . . . . . . . . . . [O III] 4959 . . . . . . . . . . . . . . . . . . . . . . . . [O III] 5007 . . . . . . . . . . . . . . . . . . . . . . . . Blend 5538 . . . . . . . . . . . . . . . . . . . . . . . . . [N II] 5755 (blue) . . . . . . . . . . . . . . . . . . [N II] 5755 (red) . . . . . . . . . . . . . . . . . . . He I 5876 (blue) . . . . . . . . . . . . . . . . . . . . He I 5876 (red) . . . . . . . . . . . . . . . . . . . . . [O I] 6300 . . . . . . . . . . . . . . . . . . . . . . . . . . [S III] 6312 . . . . . . . . . . . . . . . . . . . . . . . . . Si II 6347 . . . . . . . . . . . . . . . . . . . . . . . . . . . [O I] 6363 . . . . . . . . . . . . . . . . . . . . . . . . . . Si II 6371 . . . . . . . . . . . . . . . . . . . . . . . . . . . Blend 6548 ] 6563 ] 6584 . . . . . . [N II] 6548 . . . . . . . . . . . . . . . . . . . . . . . . . H I 6563 . . . . . . . . . . . . . . . . . . . . . . . . . . . . [N II] 6584 . . . . . . . . . . . . . . . . . . . . . . . . . Blend 6717 ] 6731 . . . . . . . . . . . . . . . [S II] 6717 . . . . . . . . . . . . . . . . . . . . . . . . . . [S II] 6731 . . . . . . . . . . . . . . . . . . . . . . . . . . He I 7065 . . . . . . . . . . . . . . . . . . . . . . . . . . . Mean . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Standard deviation . . . . . . . . . . . . .
... ... ... ... ... ... ... ... ... 1.09 ... 0.96 ... 1.00 1.03 1.03 0.99 1.02 1.03 ... 0.92 ... 1.08 ... 0.98 0.97 ... 1.01 0.05
0.69 1.14 0.74 0.71 1.01 0.92 1.04 0.90 0.90 ... 1.26 1.23 1.03 1.01 ... ... ... ... ... 1.03 ... 1.00 0.97 0.92 0.85 0.98 1.04 0.97 0.15
1.14 1.12 0.84 0.84 1.42 0.98 1.17 0.95 0.95 ... 1.27 1.24 1.07 1.05 ... ... ... ... ... 1.05 ... 1.07 1.19 0.96 1.00 0.93 1.01 1.06 0.14
1.76 1.19 0.99 0.96 1.34 1.12 1.36 1.00 1.00 ... 1.74 1.71 1.16 1.13 ... ... ... ... ... 1.26 ... 1.32 1.47 1.38 1.52 1.61 0.96 1.30 0.27
the line which BFM called C III ?4652, which actually is a blend of many weak lines). Table 2 also includes a comparison to the low-resolution results at BFM positions 1 and 3, which lie 8A east and west of position 2, respectively, in order to give an idea of the di†erences that can be expected from small positioning changes. As well as di†erent centering, we also have di†erent orientation of the slit along which
FIG. 5.ÈLog of Fe0 to Fe3` ionization fractions vs. depth in cloud
the spectra were extracted : east-west in case of BFM and at di†erent parallactic angles here. In addition to measurement errors, real di†erences in the spectrum may be important. These di†erences can be in the intrinsic spectrum (e.g., higher ionization) or they can be the result of di†erent foreground reddening (BFM). We conclude that because of systematic errors in the echelle Ñux calibration, the relative strengths of individual
FIG. 6.ÈH I velocities in the Orion Nebula with respect to the Orion rest frame (deÐned by H I recombination lines).
242
BALDWIN ET AL.
intermediate and strong lines might generally be accurate to only ^15%È20%. The Ñux calibration potentially has an additional problem near the positions of the Balmer lines because the wings of the hydrogen lines in the standard stars cover almost half of an echelle order. This is why we use He I j6678 as a normalizing line instead of Hb. Table 1 shows that the echelle measurements of the lower Balmer lines (Ha, Hb, Hc) have the same sort of scatter compared with the RC spectrograph values as do the other lines, so this problem does not appear to be especially serious for wellseparated Balmer lines. However, near the Balmer jump, the lines in the standard star spectra overlap. Comparing the intensity ratios of the higher Balmer lines as given in Table 1 to the predictions of Hummer & Storey (1987), it is quite clear that the echelle order that includes lines higher than H15 (3711 A ) and the Balmer jump does have signiÐcant Ñux calibration problems. 4.
Vol. 129
FIG. 7.ÈHe I velocities in the Orion Nebula
LINE IDENTIFICATIONS
Table 1 lists 444 emission lines measured in the 3498È 7468 A range for the combined sets of observations. Of these, we were able to identify 411 lines. From them, 37 lines are seen in both blue and red spectra. Therefore, the total number of distinct lines is 374. A number of lines are blends or have several possible identiÐcations. To identify these lines, we used atomic data from compilations by Moore (1972), Wiese et al. (1966, 1969, 1996), Fuhr, Martin, & Wiese (1988), Verner, Verner, & Ferland (1996) ; on-line databases by P. van Hoof,7 D. A. Verner,8 and NIST9 ; and earlier identiÐcations of lines in the Orion nebula made by Kaler et al. (1965), OTV, and Esteban et al. (1998). We also used Fe II energy level data (Johansson 1978 and private communication) to identify Fe II lines. We assumed the standard chemical abundance in the Orion Nebula and searched for lines from the most abundant elements Ðrst. We also assumed that the calculated velocities for lines from the same multiplet are within the errors of our measurements and within 15 km s~1 of one another. To make line identiÐcations, we paid special attention to branching ratios in multiplets and checked if other lines, especially the strongest ones, were observed. The search for other possible theoretical lines from the same multiplet allowed us to avoid some misidentiÐcations. The table lists only detections, but we also set a 3.5 p upper limit of I(He II 4686)/I(He I 6678) \ 0.0012. This limit will be important for probing the stellar continuum at ionizing wavelengths, as described below. The red observations were made purposely at a time when the EarthÏs orbital velocity caused a maximum separation between the night sky lines and the same features from Orion. The [O I] j5577 line is clearly separated from the night sky feature (by ]49 km s~1) and is detected with reasonable S/N (Table 1 and Fig. 3). The new observations give I(6300 ] 6363)/I(5577) \ 77 ^ 12. This compares very favorably with the ratio of 98 predicted by B96 using the BFM model parameters. The di†erence can be assigned to the combination of observational errors and reddening corrections and the fact that 7 Available at http ://www.pa.uky.edu/8 peter/atomic/. 8 Available at http ://www.pa.uky.edu/8 verner/atom.html. 9 Available at http ://physics.nist.gov/cgi-bin/AtData/main–asd.
the BFM model did not try to reproduce this ratio. This line was independently observed by Esteban et al. (1999) at the di†erent position, with the nondereddened ratio I(6300 ] 6363)/I(5577) \ 89 (compare to our nondereddened ratio equal to 88). The detection of the j5577 line strongly rules out a high-density region as the origin of the low-ionization lines, as shown by B96. 5.
VELOCITY FIELD ANALYSIS
Emission lines from ions with di†erent ionization potentials originate at di†erent positions along the line of sight from the H II region to the photodissociation region (PDR). As an example, Figures 4 and 5 show the ionization structure expected for O IÈO III and Fe IÈFe IV in a typical model of the Orion Nebula (B96) calculated using the development version of the photoionization code Cloudy (Ferland et al. 1998).10 The goal of this section is to use the velocity structure of the nebula to identify experimentally which lines come from which ionization zones.
FIG. 8.È[Fe II] velocities in the Orion Nebula
10 The Cloudy code and its documentation are available at http :// www.pa.uky.edu/8 gary/cloudy.
No. 1, 2000
FAINT EMISSION LINES IN THE ORION NEBULA
FIG. 9.È[Fe III] velocities in the Orion Nebula
Previous work by Balick et al. (1974) shows, from the dependence of velocities on ionization potential for a few species, that the gas is Ñowing o† the background PDR and accelerating toward us. High-ionization lines ([S IV], [O III], and [Ne III]) have negative velocities from about [2 to [1 km s~1 (blueshifted relative to the Orion rest frame) ; intermediate velocities are observed for intermediate-ionization lines ([N II] and [O II]), and positive velocities of about ]9 km s~1 are found for low-ionization lines ([Fe II] and [C II]). Our atlas adds many new lines and additional ions to this picture. There are a few ions with a large number of emission lines. There are 43 [Fe II] emission lines in both spectra, 41 lines of H I, 39 lines of He I, 32 lines of N II, 56 lines of O II, and 17 [Fe III] lines. This increases the statistics and allows us to look for correlations between velocity and excitation potential. Recombination lines should form across much of the H II region. Velocities of most H I lines (Fig. 6) have values of 0 ^ 2 km s~1. A few are in the range [6 ^ 2 km s~1, but these are the four highest Balmer lines where blending is most severe, plus 83721.839, which has twice the FWHM of the other H I lines and is therefore also a blend. All He I lines (Fig. 7) are within 0 ^ 3 km s~1.
FIG. 10.È[Fe III] velocities vs. excitation energy of the upper level
243
FIG. 11.ÈVelocities of forbidden lines vs. the ionization potential. Errors are standard deviations of the distribution of velocities from several lines for a given ion. Other ions such as [Ne III], [S III], [Ar III], and [Ar IV] are not shown because the rest wavelengths of their lines are less accurately known. Generally, they show the same trend.
Forbidden collisionally excited lines, permitted recombination lines, and pumped lines of the same ion originate from di†erent and distinct regions of the cloud. For example, [O II] lines arise in the O` zone, O II lines formed by recombination arise in the O2` zone, and O II lines pumped from the ground state arise in the O` zone. The recombination contribution to [O II] lines arises in the O2` zone, but this contribution is very small. For the neutral lines of interest, collisional contributions should be small since the electron density and temperature fall dramatically beyond the H I ionization front. Most of the large atomic O region in Figure 4 produces few strong lines in the optical band. In the following subsections we will study separately forbidden and permitted lines. 5.1. Forbidden L ines Forbidden lines are collisionally excited in the Orion Nebula (OTV). The emission from di†erent ionization stages comes from di†erent parts of the Orion Nebula and depends on physical conditions (temperature and density, which are functions of distance from the central radiative sources). The forbidden lines of nine di†erent ions from our observations were used to investigate the dependence of velocity on ionization potential. [Fe II] lines (Fig. 8) have a range of velocities from 6 to 22 km s~1, with most of them between 10 and 15 km s~1. The observed [Fe III] lines (Fig. 9) lie between 4000 to 5600 A and show a systematic increase in their velocities with their wavelengths. No other ion shows this, so it is not an instrumental e†ect. Figure 10 plots the [Fe III] velocity against the excitation potential of the upper level of the corresponding line. We see a deÐnite trend unique to this ion : the higher the excitation potential, the lower the velocity. We interpret this as an indication that there is an excitation gradient across the [Fe III] emitting region and that the gas receives the majority of its acceleration across this region. Fe2` exists over a narrow range in depth, just in front of the H I ionization front (see Fig. 5). Only this ion exhibits a velocity trend because, apparently, it happens to form across the right region where the velocity gradient is
244
BALDWIN ET AL.
Vol. 129
coincide exactly. We looked for but did not Ðnd a correlation between excitation potential and velocity in the [O II] spectrum. This is surprising since the range of excitation potentials is greater for the [O II] lines than for the [Fe III] lines. Excitation processes other than collisions can enter for the [O II] spectrum. Examples include photoionization into excited states or cascades from radiatively excited levels. This will be a focus of future work. Figure 11 plots the average velocities of various forbidden lines against the ionization potential of the corresponding parent species. Two distinct groups are obvious. The Ðrst group has velocities from ]10 to ]15 km s~1 and has lines of ions with ionization potentials less than 20 eV, namely [O I], [N I], [Fe II], and [Ni II]. The second group includes lines of ions with ionization potential larger than FIG. 12.ÈN II velocities in the Orion Nebula
large. Also, only the Fe2` ion has a large series of the collisionally excited forbidden lines with the lower level in the ground term. Figures 4 and 5 show the ionization structures of oxygen and iron. It turns out that the O` and Fe2` regions nearly
FIG. 13.ÈO II velocities in the Orion Nebula
FIG. 14.ÈVelocities of all emission lines observed in the Orion Nebula. The arrow shows the CO velocity (Balick et al. 1974). The middle line is a median, and boxes and error bars are 10th, 25th, 75th, and 90th percentiles. Velocities outside of the 10th and 90th percentiles are shown by dots.
FIG. 15.ÈComputed and observed line ratios for three stellar continua. Kurucz (1991) continua are not shown since these are known to be too soft, as evidenced by the too-small [Ne III] intensities. Horizontal bars show the dereddened line ratios for [O III] j5007 and [Ne III] j3869, and the limit to He II j4686, all relative to Hb.
No. 1, 2000
FAINT EMISSION LINES IN THE ORION NEBULA
20 eV ([S II], [O II], [N II], and [Fe III]) and has velocities around ]3 km s~1, with possibly a slight trend of decreasing velocity with increasing ionization potential (although the trend is smaller than the error bars). The similar ionization potentials of the parent ions suggest that [Ni II], [Fe II], [N I], and [O I] should all form in the same region, and this is conÐrmed by the fact that lines of these ions have similar velocities. Therefore, the average conditions where the [Fe II] originates should be close to the density and temperature determined from intensity ratios of the [N I] and the [O I] lines. 5.2. Permitted L ines Velocities of all lines with preliminary identiÐcations as N II and O II lines are plotted on Figures 12 and 13, respectively. Most of the permitted N II and O II lines are observed at similar velocities, about 0 to [5 km s~1. In a few cases, the measured velocities are quite di†erent from this average. Besides the possibility that these lines are misidentiÐed or blended, there may be physical explanations for the discrepancies. First, the lines may have a di†erent origin (recombination vs. radiatively pumped lines). Second, these contributing lines may come from di†erent parts of the Orion Nebula. This should be taken into account if these lines are used for temperature and density diagnostics. Figure 14 presents in pictorial form all the velocity information. All elements are placed in order of increasing atomic number. The middle line for each symbol is a median, and boxes and error bars are 10th, 25th, 75th, and 90th percentiles. Velocities outside the 10th and 90th percentiles are shown by dots. However, we excluded from Figure 14 lines with very discrepant velocities. We also show by the arrow the velocity of CO molecule taken from Balick et al. (1974). The CO emission comes from the background giant molecular cloud. Figure 14 summarizes the preceding discussion : 1. Most of the observable ions have velocities in the range 0 ^ 5 km s~1 (lines of H I, He I, N II, [O II], [O III], O II, [Ne III], [S II], [Si II], [Si III], [Ar III] and [Ar IV], [Fe III]). 2. Emission from another set of ions, which includes lines of [N I], N I, [O I], [Fe II] and [Ni II] mainly arises in the velocity range around 13 ^ 3 km s~1. The CO velocity of the background cloud (Balick et al. 1974) is close to this value. 3. For the same chemical element, the higher ionization stages have the smaller velocity : [N I], N I, [N II], N II, N III, O II, [O I], [O II], and [O III] ; Si II, Si III, and Si IV ; and S II, S III, and [S II], [S III], and [Fe II], [Fe III]. The [Ar III] and [Ar IV] spectra show the opposite behavior, but this result is based on very poor statistics. This is consistent with a dynamic model in which lines of ions with di†erent ionization potentials originate in gas Ñowing away from the ionization front toward the ionizing stars (and the observer too in this case). The lines from the species that are most ionized originate in the most rapidly moving (relative to the ionization front) gas closest to the stars. 6.
IMPLICATIONS FOR THE NEBULAR ENVIRONMENT
6.1. T he Acceleration across the Fe III Region The [Fe III] spectrum is unique because of the observed correlation between velocity and excitation potential (Figs.
245
9 and 10). Our interpretation is that the ionization structure of Fe2` is such that the line forms across the region where the gas receives the greatest acceleration as it moves out of the low-pressure PDR into the higher pressure H II region. Figure 5 shows the computed ionization structure of iron. Fe2` is present across a region with a thickness of 2 ] 1016 cm, and the spectroscopy shows that the radial velocity decreases by 10 km s~1. The acceleration needed to produce the observed change in velocity over the inferred distance is a \ v2/2x B 2.5 ] 10~5 cm s~2. This is an important test of hydrodynamical models of realistic H II regions and the Ðrst direct measurement of a pressure-induced acceleration within the ionized gas. The observation of an excitation potential-velocity correlation requires that there be a correlation between the temperature and velocity gradients. Photoionization models predict that this region should be nearly isothermal. Our calculations, like all simulations of the emission from H II regions, neglect hydrodynamical e†ects. In particular, the e†ects of advection are not included. This could signiÐcantly alter the thermal structure and will be a focus of future development. 6.2. Implications for the Stellar Continuum The emission-line spectrum has great sensitivity to the form of the stellar continuum. Mathis (1982, 1985) and Rubin et al. (1991) document one example in Orion, the [Ne III] j3869 line. The ion has the largest ionization potential of the observed forbidden lines, and its intensity is very sensitive to the stellar continuum near 3.5 ryd. Rubin et al. had to modify (increase) the high-energy portion of a Kurucz Atlas atmosphere to account for the intensity of this line in their studies of Orion. Sellmaier et al. (1996) show that the new generation of stellar atmospheres does produce more radiation in this critical region. The analysis presented below suggests that they may in fact produce too much. Our very deep spectrum, and the absence of He II 4686 emission, extends the possible tests of the stellar continuum. We do not detect He II j4686, and the observed limit of He II j4686/He I j6678 \ 0.0012 corresponds to a dereddened limit of less than 0.00174. This in turn corresponds to the limit He II j4686/Hb \ 7 ] 10~5. Figure 15 presents line ratios from a series of photoionization calculations that are closely patterned after those given in Baldwin et al. (1991) and Baldwin et al. (1996). The Ñux, density, and chemical composition were identical to that given in the 1996 paper, but, for simplicity, we assume constant density rather than the hydrostatic isobaric case. These parameters are close to those known to successfully reproduce the global spectrum. Only the stellar temperature and the type of atmosphere model were changed. Figure 15 shows the result of varying the temperature between 30,000 and 45,000 K for a simple blackbody and for CoStar (Schaerer & de Koter 1997) and Mihalas (1972) model atmospheres. We made similar computations with the Kurucz (1991) grid, but the results are not shown since these continua are known to be too soft and also do not produce He` ionizing photons over this temperature range. The Stoy (1933) method of determining stellar temperatures measures the shape of the ionizing continuum (Kaler 1978). Basically, the intensity of a strong cooling line relative to a recombination line is proportional to the cooling per recombination, which is equal to the heating per photoionization, which, in turn, is correlated with the hard-
246
BALDWIN ET AL.
ness of the radiation Ðeld. This ratio measures the continuum at wavelengths just shortward of 912 A . The midpanel of Figure 15 shows the [O III] j5007/Hb ratio, which closely tracks the Stoy method (Kaler 1978), and shows that stellar temperatures between 36,000 K and 37,000 K are suggested. These numbers are close to standard values (Osterbrock 1989). The [Ne III]/Hb ratio was discussed in the papers by Rubin and Mathis and is very sensitive to the continuum with wavelengths just longward of the He` ionization limit at 4 ryd. The ionization potential of Ne2` is signiÐcantly more energetic than the continuum sensed by the [O III] Stoy ratio shown in the midpanel. The CoStar atmospheres are signiÐcantly too bright, the black body too soft, while the Mihalas atmosphere works well. Our limit to the He II/Hb ratio puts simple photoncounting limits on the intensity of the continuum with wavelengths shortward of the He` Lyman limit (Osterbrock 1989). The blackbody and CoStar are far too strong here, with the latter too bright by an order of magnitude. Clearly even deeper spectra, which actually detected nebular He II j4686, would provide powerful constraints on the emergent continuum from the Trapezium stars. The surprising result is that the Mihalas (1972) atmospheres are the most successful in reproducing the spectrum of the Orion Nebula, while the more recent CoStar models do signiÐcantly worse than a simple blackbody. The He II j4686/Hb ratio is almost completely independent of the ionization parameter, assumed geometry, gas density, or abundance. This is because each photon capable of fully ionizing helium is absorbed by helium, and eventually produces a He II recombination line (Osterbrock 1989, chap. 2). The ratio is a direct measure of the number
of photons with energies greater than 4 ryd relative to the number with energies greater than 1 ryd. Numerical simulations show that this simple photon counting works over a wide range of conditions and geometries. 7.
SUMMARY
We presented the most detailed atlas to date of spectral lines in the Orion Nebula in the wavelength range from 3498 to 7468 A . These data can be obtained in electronic form11 and will be used in later papers that center on models of the formation of the spectrum. The measured intensity of [O I] j5577 conÐrms our previous limit on the electron density in neutral regions, in agreement with the Esteban et al. (1999) result. This shows that the [O I] forms at nebular densities and temperatures. Based on the analysis of the velocity Ðeld, we showed that the forbidden lines fall into two distinct groups, with implications for the hydrodynamics of the environment. The [Fe III] spectrum actually measures the acceleration of the gas across the region just this side of the H II ionization front. Finally, the Mihalas (1972) atmospheres are the most successful in reproducing the high-ionization potential spectrum. Research in Nebular Astrophysics at the University of Kentucky is supported by NSF through AST 96-17083, NASA through NAG5-4235 and 5-6937, and by STScI through AR 08387. Research by P. G. M. is supported by National Science and Engineering Research Council of Canada. R. H. R. is supported by HST GO 7514. 11 Available at http ://www.pa.uky.edu/Dkatya/orion1999.
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