The emission line spectrum of the UV deficient quasar Ton 34: evidence of shock excitation?
aa r X i v : . [ a s t r o - ph ] O c t Astronomy & Astrophysics manuscript no. krongold c (cid:13)
ESO 2018December 30, 2018
The emission line spectrum of the UV deficient quasar Ton 34:evidence of shock excitation?
Luc Binette , and Yair Krongold D´epartement de Physique, de G´enie Physique et d’Optique, Universit´e Laval, Qu´ebec, QC, G1K 7P4 Instituto de Astronom´ıa, UNAM, Ap. 70-264, 04510 M´exico, DF, M´exicoReceived: 12th July 2007/ Accepted: 20th September 2007
ABSTRACT
Context.
Emission lines in quasars are believed to originate from a photoionized plasma. There are, however, someemission features which appear to be collisionally excited, such as the Fe ii multiplet bands. Shortward of Ly α , therealso are a few permitted lines of species from low to intermediate ionization. Aims.
Ton 34 ( z q = 1 . ν ∝ ν − . ) shortward of 1050 ˚A.This object also emits unusually strong low to intermediate excitation permitted lines shortward of the Lyman limit. Methods.
Using archive spectra of Ton 34 from HST, IUE and Palomar, we measure the fluxes of all the lines presentin the spectra and compare their relative intensities with those observed in composite quasar spectra.
Results.
Our analysis reveals unusual strengths with respect to Ly α of the following low to intermediate excitationpermitted lines: O ii +O iii (835 ˚A), N iii +O iii (686–703 ˚A) and N iii +N iv (765 ˚A). We compare the observed linespectrum with both photoionization and shock models. Conclusions.
Photoionization cannot reproduce the strengths of these far-UV lines. Shocks with V s ≃
100 km s − turnout to be extremely efficient emitters of these lines and are favored as excitation mechanism. Key words.
Line: identification — Line: formation — Atomic processes— Galaxies: quasars: emission lines — quasars:individual: Ton 34
1. Introduction
In this work, we analyze the emission lines of an unusualquasar, Ton 34, which is alternatively named PG 1017+280or J1019+2745 with redshift z q = 1 . sed ) shows a remarkable steepening of thecontinuum in the rest-frame far-UV, shortward of 1100 ˚A(Binette & Krongold 2007, hereafter BK07; Binette et al.2007). If the far-UV is fitted by a powerlaw (F ν ∝ ν + α ), theindex is as steep as ν − . . BK07 suggest that the extreme-UV flux might undergo a recovery shortward of 450 ˚A.While the near-UV emission-line spectrum appears tobe ‘normal’, the far-UV spectrum shows low to intermedi-ate ionization species with unusual strengths. Using the UV sed constructed by BK07 from archive data, we will quan-tify this statement and present photoionization and shockmodels for comparison. The aim is to understand how theextreme deficiency of ionizing photons in Ton 34 might beimpacting the emission line spectrum.The emission-line spectrum of quasar and Seyfert Igalaxies is generally believed to originate from gas pho-toionized by a nuclear UV source. State of the art pho-toionization models of the Broad Emission Line Region(BELR) such as those developed by Baldwin et al. (1995)and dubbed ‘locally optimally emitting clouds’ (LOC) mod-els can successfully reproduce most of the emission lines Among the 77 quasars whose far-UV indices could be mea-sured by Telfer et al. 2002, there were only 3 objects with acontinuum steeper than ν − . observed in quasars. A grid of such models can be found inKorista et al. (1997, hereafter KO97) and more recently inCasebeer et al. (2006 and references therein). There are,however, a few exceptions to the success of pure pho-toionization. In particular, photoionization models requiremicro-turbulences in order to reproduce the shape and in-tensity of the Fe ii UV-band (Baldwin et al. 2004). A pos-sible alternative is that the region producing Fe ii is col-lisionally ionized, as proposed by Grandi (1981, 1982),Joly (1987), V´eron-Cetty et al. (2004, 2006) and Joly et al.(2007). In this work, we present evidence that photoioniza-tion might not be sustainable in the case of some of thefar-UV permitted lines reported in this paper.
2. The UV emission line spectrum of Ton 34
Below we summarize the procedure used by BK07 to derivethe UV sed of Ton 34.
The current work is based on four archival or bibliograph-ical sources. The 760–1120 ˚A spectral segment is providedby the dataset Y2IE0A0AT from the HST-FOS archives(grating G270H). To cover the extreme UV region, weborrowed from the IUE archives. The long wavelengthsegment (LWP) is from Tripp, Bechtold & Green (1994)and corresponds to the dataset LW0P5708. Fluxes long-ward of 3000 ˚A (observer-frame) were severely affectedfrom reflected sunlight or moonlight (Lanzetta, Turnshek &
Binette & Krongold: The emission lines of Ton 34
Sandoval 1993) and have been discarded. The shorter wave-length IUE segment (SWP) was extracted directly fromthe archives and corresponds to the dataset SWP28188.To cover the sed behavior longward of the HST seg-ment, we adopted the published optical spectra of Sargent,Boksenberg & Steidel (1988), which were taken at thePalomar 5.08 m Hale Telescope. Both optical spectra lackedabsolute flux calibration, although the authors observedstandard stars, which allowed them to provide a relativecalibration. sed segments
We statistically corrected the UV spectral segments forthe cumulated absorption caused by unresolved Ly α forestlines, which are responsible for the so-called far-UV “Lymanvalley” (Møller & Jakobsen 1990). For that purpose, weadopted the mean transmission function for z q = 2 pub-lished by Zheng et al. (1997). We also applied a Galacticreddening correction assuming the Cardelli, Clayton &Mathis (1989) extinction curve corresponding to R V = 3 . B − V = 0 .
13. The latter value corresponds to the meanextinction inferred from the 100 µ maps of Schlegel et al.(1998) near Ton 34. The blue and red arm segments havebeen scaled to overlap smoothly with the HST-FOS seg-ment. Both the LWP and SWP segments were multipliedby a factor 0.75. This scaling was necessary so that the LWPsegment superimposes the HST-FOS spectrum as closely aspossible. Continuum variability is a possible explanation forthis continuum difference, since the IUE and HST observa-tions were made in different years. Finally, all the spectralsegments were shifted to rest-frame wavelengths, and F λ was multiplied by 1 + z q . The IUE spectra have been re-binned by grouping n pixels together [SWP with n = 5 andLWP with n = 3) in order to improve the limited S/N. TheLWP and HST-FOS spectra overlap significantly in spec-tral coverage. Both datasets taken nine year apart confirmthe unusual steepness of the UV break in Ton 34. sed of Ton 34 Shortward of 1100 ˚A, the continuum of Ton 34 undergoes asharp fall off (see Fig. 2 in BK07), which BK07 model asdust absorption by nanodiamond grains. This resulted ina deep and broad absorption trough that fits the observedcontinuum reasonably well. In our photoionization calcula-tions presented below in Sect. 3.2.1, we experiment with twoionizing sed s. The first is the intrinsic ‘unabsorbed’ sed ,which is assumed to be a powerlaw of index +0 . keV , sed ii behaves as a powerlaw ofindex − .
0, yielding an α OX of − sed is shownin Fig. 1 and, as in the work of BK07, it is labeled Model ii .The second sed used in photoionization calculations is thedust absorbed version of the same sed , which fits the ob-served UV continuum of Ton 34 between 400 and 1550 ˚A (la- This correction is statistical in nature, as it relies on theaverage behavior with redshift of the spatial density of inter-vening absorbers. It cannot be used to correct small portions ofthe continuum, which may be coincident with a “clear patch” oran over-density in the Ly α forest. These inhomogeneities maygenerate spurious narrow features that should not be attributedto genuine emission lines. Fig. 1.
Log-log plot of the input spectral energy distribu-tions used in our photoionization calculations discussed inSect. 3.2.1. These ionizing sed s are labeled ii and iv in ei-ther F λ (panel a ) or F ν (panel b ) and are given as a functionof wavelength (bottom axis) or photon energy (top axis).The distribution labeled ii (solid line) is the assumed in-trinsic sed while that labeled iv is the transmitted flux(dash-dotted line) assuming nanodiamond dust extinction(see BK07). Model iv is a fit of the UV continuum of Ton 34between 400 and 1550 ˚A.beled Model iv in Fig. 1). Shortward of 200 ˚A and longwardof 2000 ˚A, the two distributions are the same. This is be-cause nanodiamond dust absorbs radiation over a relativelynarrow domain as compared to other grain compositions.In Fig. 2, we present the continuum subtracted spectrumof Ton 34, that is, the residual between the observed Ton 34 sed and our continuum fit represented by Model iv . The dif-ferent spectral segments have been color-coded as follows,SWP: red, LWP orange, HST-FOS: blue, and Palomar:dark green. The procedure to measure the flux of the lines was thefollowing: we first fit a Gaussian to each observed line inthe spectra. For several lines, a narrow component was re-quired, so we added a second (narrow) Gaussian. In addi-tion, the lines by C iv λ iv λ α show aclear asymmetry in the line profile, with a blue shoulder (seeFig. 2). For these lines, we further included a third, broader,Gaussian. The FWHM of the broad component spans from ∼ − . It is interesting to note that theO ii +O iii complex at around 835 ˚A has a significant andstrong red shoulder extending up to ∼
850 ˚A, which is ob-served in both the IUE-LWP and HST-FOS spectra (seeFig. 2). We could not find any positive identification of thisshoulder with any line from a different ion/transition, and inette & Krongold: The emission lines of Ton 34 3
Fig. 2.
Residuals of the spectral energy distribution ofTon 34 after subtracting our absorbed continuum Model iv from BK07. The different spectral segments have been colorcoded as follows, SWP: red, LWP orange, HST-FOS: blue,and Palomar: dark green. Color-coded fiducial marks in-dicate the position of observed or expected (labeled withsymbol ’ ?’) emission lines. Measurements of line intensitiesand upper limits are given in Table 1.thus we considered this feature as part of the O ii +O iii emission.The measured line fluxes extracted from Fig. 2 as well asupper limits of other permitted lines are listed with respectto Ly α = 100 in Col. 5 of Table 1. Note that we give thetotal flux under the profile, that is, the integrated flux fromall the Gaussian components required to fit each emissionline. A consistency check was carried out, which showedthat the line fluxes measured over the original spectra orthe continuum subtracted spectra were indistinguishablefrom each other.In Col. 5 of Table 1, we show our error estimates, whichwe evaluated at a 1 σ significance level. We assumed a S/Nof 25 for most lines, except for N iii +N iv and N iii +O iii where we assume a S/N of ≃
10. The line upper limitsin Table 1 correspond to a significance of 2 σ . As for thecontinuum, we estimate the errors to be ≃ ii +O iii lines at 835 ˚A, the N iii +O iii lines at 686–703 ˚A and theN iii +N iv lines at 765 ˚A.Many weaker features in the IUE spectrum appear to liewhere other permitted lines of comparable excitation mightbe expected, such as O iii λ
508 ˚A, O iv λ
554 ˚A, O v λ
630 ˚Aand O iv λ
609 ˚A. A few of these have been reported beforein other quasars (Reimers et al. 1998; Laor et al. 1995) orin composite AGN spectra (Zheng et al. 1997; Telfer et al.2002; Scott et al. 2004). However, these line systems appearas too narrow in the IUE spectra with respect to typical
Fig. 3.
Emission lines extracted from the Ton 34 spectrumplotted in velocity space. The flux scale is arbitrary for eachinset. Left panels: near-UV permitted lines, right panel: far-UV permitted lines. Overall, all the lines are consistent withthe rest frame system of Ton 34. Differences in the positionof the lines on the right panel may be due to absorption byintergalactic gas. The narrow line of O iii at 508 ˚A (bottomright panel) is severely affected by intergalactic absorption,and better data would be required to confirm its presence.The same applies to the other lines shown as upper limitsin Table 1.BELR line profiles (see profile comparison of Fig. 3). Theylack a broad component at their base. Given the limitedS/N of the IUE spectrum at the far-UV end, we considerprobable that these lines are spurious features instead. Forthis reason, we will consider these emission-like features asupper limits rather than real detections. The symbol ‘?’denotes these unconfirmed lines in our various figures.We find little evidence of the high excitation Ne viii line at 775 ˚A reported by Telfer et al. (2002) and Scott et al.(2004) in their respective composite spectrum, and we favorthe identification of O iv λ
789 ˚A instead. Because the linespectrum of Ton 34 is of unusually low excitation as shownbelow in Sect. 3.1, we do not believe that the high excitationlines of Mg x and Ne viii (listed in Table 1) are present ata detectable level. As can be gathered from Fig. 2, the strongest emission fea-tures in the far-UV coincide with the position of lines ob-served or expected in quasar spectra (Sect. 3.1). However,the limited quality of the data and the possible coinci-dence of absorbers at inconvenient spectral positions pre-vent us from deriving incontrovertible conclusions. In thecase of the narrower features (O iii λ
508 ˚A, O iv λ
554 ˚A,O v λ
630 ˚A and O iv λ
609 ˚A), better quality data is re-quired to confirm or discard their presence, as discussed in
Binette & Krongold: The emission lines of Ton 34
Fig. 4.
Line flux ratios renormalized to Ly α = 100 of differ-ent species as a function of line wavelength (˚A). These wereextracted from the the spectrum of Ton 34 (filled squares)and from the radio-loud (crosses) and radio-quiet (openlozenge) composite spectra of Telfer et al. (2002). Only ra-tios that have a counterpart in Ton 34 are shown. Largersquares correspond to ratios for which the difference be-tween Ton 34 and the composites exceeds a factor 10 (alsoshown in bold face in Col. 5 of Table 1). The symbol ’ ?’ de-notes upper limits of unconfirmed lines in Ton 34. The grayfilled triangle indicates the position of Ly α .Sect. 2.4. Clearly, new observations are needed in all wavebands down to the X-rays. In what follows, we will take thedata at face value and present photoionization and shockmodels that attempt to reproduce the far-UV lines.
3. Modelling of the line spectrum
We now quantify to which degree the emission lines dif-fer in Ton 34 from the ‘average’ quasar. To achieve this, welist the line ratios characterizing the radio-loud (Col. 3) andradio-quiet (Col. 4) composite spectra of Telfer et al. (2002)in Table 1. Comparison between Ton 34 and these two setsof ratios require some caution, since significant line ratiovariations exist among quasars. For instance, Telfer et al.(2002) reported that the RMS deviation of line fluxes be-tween the different quasars amounted to as much as 50–70% for the strong lines of C iv λ vi λ α .Hence, intrinsic differences of less than a factor two be-tween the composites and Ton 34 should not be consideredsignificant.To facilitate the comparison of Ton 34 with the two com-posites, we plot their line ratios in Fig. 4. Inspection of theTable 1 or Fig. 4 reveals that the commonly strong BELRlines of C iv , N v and O vi are all present in Ton 34. Hencethe apparent sharp turndown of the ionizing UV in the range 650–912 ˚A is not affecting radically the high exci-tation emission lines. In particular, the O vi λ iv is substantially weaker, by more than a factorof six in Ton 34 with respect to the radio-quiet composite.Also, the line system C iii +N iii near 980 ˚A is noticeablyweaker, although the flux in this line is difficult to measureaccurately due to the uncertainties introduced by the sharpcontinuum bent and the many Ly α forest lines.In the far-UV, we note that the intensity of theO ii +O iii and N iii +O iii systems in Ton 34 are a factorof ∼
14 and 18 brighter, respectively, than in the RLQcomposite. There is also evidence of significant emission ofN iii and/or N iv at 764 and 765 ˚A, which are not detectedin the composite spectra either. The line spectrum of Ton 34 show peculiarities that de-serve further analysis. In particular, O ii +O iii (835 ˚A),N iii +O iii lines (686–703 ˚A) and N iii +N iv (765 ˚A), whichare measured with unusual strengths with respect to Ly α .Are these emission features necessarily genuine lines? Onepossibility is that extinction resonances, unaccounted forin the extinction curve used to model the deep continuumtrough (BK07), may induce features that looked like broademission lines. Another possibility is that Ly α absorbers atintervening redshifts might generate spurious emission fea-tures by bracketing narrow continuum regions. Althoughwe cannot rule out either possibility with the current data,both appear unlikely to us, on the ground that the strongestemission features coincide quite well with the position ofplausible atomic transitions (see Fig. 2). The two strongestline systems of O ii +O iii (835 ˚A) and N iii +O iii (686–703 ˚A) have previously been reported in the RLQ compos-ite, although at a much reduced flux level. We will thuspursue our analysis under the assumption that the observedfeatures are real and consist of low to intermediate excita-tion permitted lines . Can photoionization account for the strength of the far-UV permitted lines? We first establish a comparison withpublished BELR models, and then evaluate the impact ofa strongly absorbed ionizing continuum.Baldwin et al. (1995) showed that by integrating linefluxes over a wide range in gas density n H and impingingionizing flux ϕ H , one obtains a much improved fit to quasarline spectra. Such models were dubbed “locally optimallyemitting clouds” (LOC). Baldwin et al. (1995) also showedthat by preferentially selecting the optimal slab density andimpinging flux for each individual line, one can derive aline spectrum comparable (within a factor two) to that ofa true LOC model. To derive an approximate LOC model,we proceeded as follows. From the grid of photoionizationmodels published by Korista et al. (1997; hereafter KO97),we extracted the highest equivalent width found within theplane ϕ H vs n H , for each line of interest. The particulargrid that we selected was labeled agn4 . It assumes solarabundances and a sed that was defined by KO97, which The selected grid agn4 ∼ korista/grids/grids.htmlinette & Krongold: The emission lines of Ton 34 5 Table 1.
Comparison of Ton 34 with composite sed s and with models
Lines Observations Photoionization a Shocks a,b
Species λ (˚A) RLQ RQQ Ton 34
KO97 c sed ii d sed iv d d km s − (1) (2) (3) (4) (5) (6) (7) (8) (9)O iii < < ≤ . − . iv < < ≤ . viii < e – ? ? ? ?He i
601 – – ≤ . iv +Mg x < < < . v < < ≤ . iii +O iii e < ± . iii +N iv < < ± . viii +O iv ≤ . ii +O iii ± .
32 ? +0.5 0.04+1.4 0.04+2.4 48+23C ii +N ii ≤ . < − < − iii +N iii ± .
64 2.9+0.7 4.9+0.8 3.9+0.5 13+7.0Ly β +O vi ± .
75 1.1+20 0.37+1.5 0.36+21.5 2.5+10 − C ii f ? 0.02 0.02 4.1N ii +He ii iv < < ± .
23 ? 0.97 0.95 1.6Fe iii < − iii iii +Si iii f α g ± . v f − N v ± . ii i +Si ii ± .
24 0.07+0.03 10 − . +0.01 10 − . +0.02 10 − . +0.15C ii ± .
13 0.7 0.63 0.7 44Si iv +O iv ] 1397, 1402 8.6 11.9 9.2 ± . iv ] 1486 2.8 0.6 1.3 ± . iv ± . ii ± ? 3.0 4.0 9.5 2.1O iii ] 1664 2.3 0.7 – 7.8 8.1 6.3 0.7 a Some observational entries in Cols. 3–5 corresponds to the sum of two different lines. For the corresponding models in Cols. 6–9,we list separately each line intensity using a + symbol as separator. b Redward of 1700 ˚A (down to the infrared), the 100 km s − shock does not generate any strong lines. For completeness, theonly other lines of significant brightness are Si iii ii ii α , respectively. As for the (optical and UV) Fe ii multiplet line systems, we can’t say since they arenot considered by mappings i c. Shortward of 400 ˚A, we expect the He ii lines to be strong, with He ii
304 ˚A reaching 80% ofLy α . c A crude model that approximates the optimally locally emitting BELR model described by Baldwin et al. (1995). Each line’speak emissivity was extracted from the grid agn4 of photoionization calculations published by Korista et al. (1997) d These three models were computed with mappings i c assuming an initial density n of 4 × cm − and solar metallicities.The ionization parameter is 0.04 for the two photoionized models and zero for the shock model. At these densities, the Ly α luminosities per unit area of photoionized or shocked gas are 3 . × , 4 . × and 5 . × erg cm − s − for models shownin Cols. 7, 8 and 9, respectively. These would scale approximately in proportion to n . e Measurement by one of us (YK) using the composite spectra lent by R. Telfer. f The strong neighboring lines of Ly α or O vi makes the determination of a meaningful upper limit impossible. g The Ly α flux in Ton 34 is measured to be 6 . × − erg cm − s − corresponding to an equivalent width of 57 ˚A. peaks at 22 eV. It is the closest to our sed ii with a 18.5 eVturnover (Fig. 1; see also Haro-Corzo et al. 2007).The line ratios from this approximated LOC model areshown in Col. 6 of Table 1. Unfortunately, the N iii +O iii line system ( λλ agn4 gridnor the O ii λ
834 ˚A line. On the other hand, the N iii +N iv system at 765 ˚A and the O iii line at 835 ˚A were. TheN iii +N iv system is a factor of a few weaker than observedwhile the λ
835 ˚A O iii line is predicted an order of mag-nitude weaker than the observed O ii +O iii system. As we consider unlikely that the λ
834 ˚A O ii line (absent from the agn4 grid) is stronger than O iii , we conclude that pho-toionization would have great difficulties in fitting this sys-tem. Hence, even locally optimally emitting clouds wouldnot be able to account for the intensities of at least someof the far-UV lines observed in Ton 34.Could the peculiar shape of the Ton 34 sed be re-sponsible for the unusual strengths of some far-UV lines?Out of curiosity, we calculated with the multipurpose code mappings i c (Ferruit et al. 1997; Binette et al. 1989) pho- Binette & Krongold: The emission lines of Ton 34 toionization models using sed ii to compare with the ab-sorbed sed iv , characterized by the deep trough. We as-sumed solar metallicities (Anders & Grevesse 1989) and agas density of 4 × cm − . The ionization parameter wasvaried until a maximum in the O iii ]/H β ( λ λ U = 0 .
04. The modelswere truncated at a depth where H is 10% ionized. Thesecalculations with U = 0 .
04 using either sed ii or iv (bothplotted in Fig. 1) are reported in Cols. 7 and 8 of Table 1, re-spectively. Because there are fewer soft ionizing photons in sed iv , we find that the mean energy of the photoelectronsis twice as high as the one given by sed ii . Hence, this mustresult in a hotter plasma and therefore in stronger collision-ally excited lines. A comparison of the calculated ratiosbetween the two models and with Ton 34 (Col. 5) revealsthat, although many metal lines in Col. 8 ( sed iv ) are oftenstronger by a factor of a few with respect to Col. 7 ( sed ii ),the deep UV trough does not result in a sufficient increasein the strengths of the O iii +N iii lines at 683, 703 ˚A nor ofthe O ii +O iii lines at 835 ˚A. In conclusion, photoionizationpredicts far-UV line intensities much too weak in compar-ison with our measurements. Furthermore, making drasticchanges in the shape of the ionizing continuum does notalter this conclusion. In view of the difficulties of producing strong permittedlines of O ii , O iii and N iii in the case of pure photoioniza-tion, we are lead to consider whether collisional ionizationmight not be more appropriate.To investigate this possibility, we calculated with mappings i c a sequence of steady-state plane-parallelshock models with a preshock density of 4 × cm − , as-suming again solar metallicities. The postshock tempera-tures of the different models covered the range 1 . × – 8 × K, corresponding to shock velocities of 75 to235 km s − . The pre-ionization state of the shocked gas wasdetermined self-consistently by an iterative scheme, usingthe ionizing radiation produced within the cooling shockthat propagates upstream (Dopita, Binette & Tuohy 1984).The time evolution of the electron and ion temperatureswere followed separately until they equalized, making useof the equilibration timescale as defined by Spitzer (1962).Most of the far-UV resonance lines are emitted downstreamin layers of densities in the range 10 . –10 . cm − , wellbelow the densities of 10 where collisional de-excitationwould become a concern for many resonance lines. Theelapsed time for the shocked gas to cool to temperaturesof 8500 K is about 10 seconds. The adiabatic cooling andrecombination of the plasma was followed in time untilthe ionized fraction reached ≤ iii λ
977 and C iv λ − shock, compared to 10 . and 10 . for the pho-toionization model of Col. 8. We use the customary definition of the ionization parameter U = ϕ H /cn H , which is the ratio of the density of ionizing pho-tons impinging on the slab ϕ H /c to the H density at the face ofthe slab n . Fig. 5.
Line intensities from high density cooling shocksrenormalized to Ly α = 100 as a function of shock velocity.Solar metallicities have been assumed. A vertical dashedline denotes the velocity of the shock model reproduced inTable 1.The intensities of representative far-UV lines are shownin Fig. 5 as a function of shock velocity. The calcula-tions show that shocks with gas densities appropriate tothe BELR are very efficient in producing strong lines ofO ii +O iii ( λ
835 ˚A) and of N iii +O iii ( λλ α , respectively. We also com-puted the intensities of many other far-UV lines that mightbe observable in future observations. Some high excitationlines such as O iv λ
554 ˚A, O iv λ
789 ˚A and O v λ
630 ˚A be-come intense for shock velocities exceeding 120 km s − . Bycomparing in Table 1 the observed upper limits for theselines with the computed intensities of O iii λ
835 ˚A or O iii λ
703 ˚A, we find that velocities of order 90–130 km s − pro-duce line intensities compatible with the estimated line ra-tios . To be definite, we adopt the velocity of 100 km s − for the case model presented in Col. 9 of Table 1.Shock models by themselves predict far-UV line inten-sities that are too strong with respect to Ly α (compareCols. 9 and 5), creating a reverse situation to that of pho-toionization (Sect. 3.2.1). We are therefore lead to proposea mixed model, in which we ascribe only a fraction of theluminosity of Ly α to be due to shock excitation and thecomplementary fraction to photoionization. In this mixedmodel, photoionization would be responsible for the emis- While the measurements for O iv λ
554 ˚A and O v λ
630 ˚Aformally represent only upper limits, it remains possible thatthe intensities of these lines are somewhat larger than evaluatedgiven the limited S/N of the IUE-SWP spectrum and the possi-ble presence of many intergalactic absorption lines (this wouldimply higher shock velocities). For completeness, Table 1 includes all the lines that theshock model predicts to be stronger than 2% of Ly α withinthe reported domain of 400–1700 ˚A.inette & Krongold: The emission lines of Ton 34 7 sion of the strong near-UV (i.e. classical) lines, while shockswould be contributing of order a third of Ly α and (propor-tionally) all of the far-UV resonance lines shortward of theLyman limit.The preshock density n may be significantly higherthan assumed above. We find similar line ratios for preshockdensities up to 100 times higher. The luminosity per unitarea of the shock model in this case exceeds that of thephotoionization models presented in Col. 7 and 8 (see foot-note a in Table 1). Our code includes three body recom-bination of H, but not the process of stimulated emission,which prevents us from going beyond a preshock densityof 10 . cm − . Beyond this limit, we expect Ly α to be thefirst line to thermalize, which would further enhance thestrengths of the metal lines with respect to Ly α .In summary, the far-UV lines observed in Ton 34 short-ward of the Lyman limit are characterized by a much higherexcitation energy than the near-UV lines. For this reason,collisional excitation (through shocks) at temperatures sig-nificantly higher than that typically provided by photoion-ization is strongly favored. Calculations with mappings i cshow that such temperature regime is ensured when shockexcitation of moderate V s takes place. These shocks wouldaccount not only for the far-UV lines, but may contributesignificantly to the Fe ii multiplet lines that have been pro-posed to result from mechanical heating by Joly et al. (2007and references therein). Acknowledgements.
This work was supported by the CONACyTgrants J-50296 and J-49594, and the UNAM PAPIIT grant IN118905.Diethild Starkmeth helped us with proofreading.
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