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In addition in the boundary. laver region, the radial velocity is large and (the accreting matter can also heat up by shock. | In addition in the boundary layer region, the radial velocity is large and the accreting matter can also heat up by shock. |
For ese reasons (he soft accretion assumption is actually not justified. since this advected energv is non-negligible. | For these reasons the 'soft' accretion assumption is actually not justified, since this advected energy is non-negligible. |
ILowever. in our treatment of the boundary laver irradiation. we lake into account part of this energy. by assuming that a fraction of the energy liberated in 1e boundary laver (L5) is absorbed by the outer laver of the star: it is included as a source in the energv equation in the outer laver of the star. | However, in our treatment of the boundary layer irradiation, we take into account part of this energy by assuming that a fraction of the energy liberated in the boundary layer $L_{BL}$ ) is absorbed by the outer layer of the star: it is included as a source in the energy equation in the outer layer of the star. |
The treatment of the boundary laver irradiation is done as follows. | The treatment of the boundary layer irradiation is done as follows. |
The energy liberated in the boundary laver is given by (μιας1987): where is the total accretion energy. G is the gravitational constant. A, is the mass of the star. R, is the radius ofthe star. AV is the mass accretion rate. O, is the angular rotation rate of the star and Oj(H,) is the IXeplerian angular velocity at one stellar radius. | The energy liberated in the boundary layer is given by \citep{klu87}: where is the total accretion energy, $G$ is the gravitational constant, $M_{*}$ is the mass of the star, $R_{*}$ is the radius ofthe star, $\dot{M}$ is the mass accretion rate, $\Omega_{*}$ is the angular rotation rate of the star and $\Omega_K(R_{*})$ is the Keplerian angular velocity at one stellar radius. |
Equation (1) can be used as long as the disk is geometrically thin and optically (hick and extends to the stellar surface. | Equation (1) can be used as long as the disk is geometrically thin and optically thick and extends to the stellar surface. |
In (he present case one expects the stellar rotation rate to be the rotation rate of the WD: O,=QO,. | In the present case one expects the stellar rotation rate to be the rotation rate of the WD: $\Omega_{*}= \Omega_{wd}$. |
We assumed (hat only a fraction of the boundary laver Iuminositv is irradiating the star. namely: A value a=1 means that half of the BL Inminosity is lost into space while the other half is absorbed by the star. | We assumed that only a fraction of the boundary layer luminosity is irradiating the star, namely: A value $\alpha=1$ means that half of the BL luminosity is lost into space while the other half is absorbed by the star. |
Assuming a value a=0.5 means thal only of the BL luminosity is absorbed by the star. | Assuming a value $\alpha=0.5$ means that only of the BL luminosity is absorbed by the star. |
Here. we choose a= 0.5. which is the value used in the | Here, we choose $\alpha=0.5$ , which is the value used in the |
Con A was observed with the CANCAROO 3.81 TeV sav telescope from March to April 1995. but was not detected (Rowell et al. | Cen A was observed with the CANGAROO 3.8m TeV $\gamma$ -ray telescope from March to April 1995, but was not detected (Rowell et al. |
1999). | 1999). |
There was no N-rav observation for Cen A in 1995. but using the observations before 1995 we can inter that Cen A was in the low state during 1995. | There was no X-ray observation for Cen A in 1995, but using the observations before 1995 we can infer that Cen A was in the low state during 1995. |
Between 1992-1995 Con A was very faint aud there was no trend towards outburst in 1995 (sec Fie. | Between 1992-1995 Cen A was very faint and there was no trend towards outburst in 1995 (see Fig. |
5 in Boud et al. | 5 in Bond et al. |
1996). | 1996). |
Thus. the TeV fiux of Cen A was below the detection limit of the CANCAROO telescope. | Thus, the TeV flux of Cen A was below the detection limit of the CANGAROO telescope. |
Crindlay et al. ( | Grindlay et al. ( |
1975) reported that Cen A was detected in the TeV cnerev range divine 1972-1971. | 1975) reported that Cen A was detected in the TeV energy range during 1972-1974. |
Since the detection has not been repeated (Ixmael 1998). it appears that no one has iucluded Con A in the list of known TeV - ACNs, | Since the detection has not been repeated (Israel 1998), it appears that no one has included Cen A in the list of known TeV $\gamma$ -ray AGNs. |
However. from 1972 to 1976. Cen A uuderwenut a large outburst in the N-vav energy range (sec Fie. | However, from 1972 to 1976, Cen A underwent a large outburst in the X-ray energy range (see Fig. |
bin Turner et al. | 4 in Turner et al. |
1997 or Fig. | 1997 or Fig. |
5 iu Boud et al. | 5 in Bond et al. |
1996). aud probably uuderweut an outburst iu the TeV energy rauge at the same time. so the detection of Cen A by Crindlay et al. ( | 1996), and probably underwent an outburst in the TeV energy range at the same time, so the detection of Cen A by Grindlay et al. ( |
1975) is considered to be authentic. | 1975) is considered to be authentic. |
At present. the CANCAROO I0 telescope is the ouly TeV οταν telescopes which cau be used to observe the southern object Cen A. Iu the future. Con A can be observed by CANCAROO III and the Ceriizur-Freuch- experiment TESS in Namibia. | At present, the CANGAROO 10m telescope is the only TeV $\gamma$ -ray telescopes which can be used to observe the southern object Cen A. In the future, Cen A can be observed by CANGAROO III and the German-French-Italian experiment HESS in Namibia. |
The loue-teru: rav light curve of Con A shows that the period of outburst ish6 vears(see Fig. | The long-term X-ray light curve of Cen A shows that the period of outburst is $5 - 6$ years (see Fig. |
5 iu Boud et al. | 5 in Bond et al. |
1996). which is typical for radio-loud ACNs(σημ and Nair 1995). | 1996), which is typical for radio-loud AGNs (Smith and Nair 1995). |
The recent observation bv the Chandra N-rav Observatory i 1999 shows that Cen A is in the low state at present (Ixraft et al. | The recent observation by the Chandra X-ray Observatory in 1999 shows that Cen A is in the low state at present (Kraft et al. |
2000). | 2000). |
Thus. it is expected that Cen A may undergo an outburst in the near future. | Thus, it is expected that Cen A may undergo an outburst in the near future. |
The long-term Πο curve also shows that Cen A did not undergo anv laree outburst since 1986. so the comiug outburst may be a large ouc. | The long-term light curve also shows that Cen A did not undergo any large outburst since 1986, so the coming outburst may be a large one. |
Iu addition. according to Eq. ( | In addition, according to Eq. ( |
1) the Compton component of Cen A av also peak as high as —100 TeV. It may thus be strong enough to be detectable even iu the PeV energy range during outburst. | 1) the Compton component of Cen A may also peak as high as $\sim$ 100 TeV. It may thus be strong enough to be detectable even in the PeV energy range during outburst. |
TheFEST has recorded that the fiux from the AIST optical uucleus varies by a factor of —2 on timescales of ~2.5 mouths (Tsvetanov ct al. | The has recorded that the flux from the M87 optical nucleus varies by a factor of $\sim$ 2 on timescales of $\sim$ 2.5 months (Tsvetanov et al. |
1998). | 1998). |
A difference bv a factor of five im the fux detected by and Cüuga indicated that the lard N-vav from M87 core also varied violeutlv (ποπα Moleudi 1999). which suggests that the variability in ALS? around peak frequency is similar to those iu Τον blazars. | A difference by a factor of five in the flux detected by and Ginga indicated that the hard X-ray from M87 core also varied violently (Guainazzi Molendi 1999), which suggests that the variability in M87 around peak frequency is similar to those in TeV blazars. |
EGRET has observed ALS? but did uot cetect auv Ce 5-ravs (Sreckumar et al. | EGRET has observed M87 but did not detect any GeV $\gamma$ -rays (Sreekumar et al. |
1996). which mav indicates that AIST does not have a peak in its Compton component at the GeV energy rauge. beiug consistent with our result. | 1996), which may indicates that M87 does not have a peak in its Compton component at the GeV energy range, being consistent with our result. |
AIST was observed at TeV energies many vears ago but was not detected. | M87 was observed at TeV energies many years ago but was not detected. |
At that time the TeV detectors were not sensitive (the CrabNebula was not detected either). aud the olservatious oulv gave upper lnuits. Ες0.UT0eV)<SS.10Hohotouscin?7s3 (Cawley ct al. | At that time the TeV detectors were not sensitive (the Crab Nebula was not detected either), and the observations only gave upper limits, $F(>0.4 {\rm TeV}) < 8.3\times10^{-11} {\rm photons}\, {\rm cm}^{-2}{\rm s}^{-1}$ (Cawley et al. |
1985) aud F(>(.21TeV)<1.2«101honscur?74 (Weekes ct al | 1985) and $F(>0.21 {\rm TeV}) < 1.2\times10^{-10} {\rm photons}\, {\rm cm}^{-2}{\rm s}^{-1}$ (Weekes et al. |
1972: see Fig. | 1972; see Fig. |
1). | 1). |
M8? has not vet been observed by auy ποσα TeV s-rav telescopes. | M87 has not yet been observed by any modern TeV $\gamma$ -ray telescopes. |
It cau be observed using powerful TeV x-ray telescopes. such as Whipple aud IIEGRÀ. | It can be observed using powerful TeV $\gamma$ -ray telescopes, such as Whipple and HEGRA. |
Iun sunny. the N-ravs from Cen A and MsST are svuchrotron enüssion rather than inverse-Coupton cluission. and hence these nearby ERI radio galaxies are IDL-like objects. | In summary, the X-rays from Cen A and M87 are synchrotron emission rather than inverse-Compton emission, and hence these nearby FRI radio galaxies are HBL-like objects. |
In particular. Con A is similar to Myk 501 with its svuchrotron compoucut peaking at very high energies. Which was ounce thought to be nou-svuchrotron radiation (Skibo et al. | In particular, Cen A is similar to Mrk 501 with its synchrotron component peaking at very high energies, which was once thought to be non-synchrotron radiation (Skibo et al. |
1991: ικα ot al. | 1994; Kinzer et al. |
1995: Steiule et al. | 1995; Steinle et al. |
1998). | 1998). |
According to the unified scheme of BL Lac objects aud FRI radio galaxies aud assmumine that the properties of the known TeV BL Lac objectsCon are common for all TeV AGNs. we predict that A ay have a peak iu its Compton component power output at ~1 TeV. and that AIST may have a Compton cussion peak at ~0.1 TeV. both having TeV οταν flux detectable by TeV οταν detectors available todas. | According to the unified scheme of BL Lac objects and FRI radio galaxies and assuming that the properties of the known TeV BL Lac objects are common for all TeV AGNs, we predict that Cen A may have a peak in its Compton component power output at $\sim$ 1 TeV, and that M87 may have a Compton emission peak at $\sim$ 0.1 TeV, both having TeV $\gamma$ -ray flux detectable by TeV $\gamma$ -ray detectors available today. |
This work was financially supported by the B21 Project of the Iorean goverinnent. | This work was financially supported by the BK21 Project of the Korean government. |
We thaws Ueluut Steiule and Tracey Jane Turucer for providing us with soft s-rav and X-ray data. respectively, | We thank Helmut Steinle and Tracey Jane Turner for providing us with soft $\gamma$ -ray and X-ray data, respectively. |
We also thank the anonviuious referee for helpful sugecstious. | We also thank the anonymous referee for helpful suggestions. |
We also thank the anonviuious referee for helpful sugecstious.O | We also thank the anonymous referee for helpful suggestions. |
We also thank the anonviuious referee for helpful sugecstious.OO | We also thank the anonymous referee for helpful suggestions. |
Unfortunately. the strengths of pronmüunent metal lines. such as CIvALS19. relative to Lya are not sensitive to the overall metallicity for Z=O12. (CIlunann Ferland 1999). | Unfortunately, the strengths of prominent metal lines, such as $\lambda 1549$, relative to $\alpha$ are not sensitive to the overall metallicity for $Z \ga 0.1 Z_\odot$ (Hamann Ferland 1999). |
Shields (1976) proposed that the relative uitrogen abundance could be used as an indirect metallicity indicator. | Shields (1976) proposed that the relative nitrogen abundance could be used as an indirect metallicity indicator. |
Asstuning that the secondary nitrogen production. ie. the svuthesis of nitrogen from existing carbon aud oxygen via CNO burning iu intermediate mass stars (Tinsley 1980: Iesv et 22000). is the dominant source fornitrogen. this results N/OxO/T aud hence ΑΠx(O/IIYiu Z. | Assuming that the secondary nitrogen production, i.e., the synthesis of nitrogen from existing carbon and oxygen via CNO burning in intermediate mass stars (Tinsley 1980; Henry et 2000), is the dominant source fornitrogen, this results in $N/O \propto O/H$ and hence $N/H \propto (O/H)^2 \propto Z^2$ . |
Observations of ΠΠ regions indicate that secondary uitrogen production aud the ΑΟx scaling dominate when the imetallieitv is above ~1/3 tfo ~1/2 sobw (Shields 1976: Pagel Edumuds 1981: van Zee et 11998: Izotov Thuan 1999: Pettini oet 22002). | Observations of HII regions indicate that secondary nitrogen production and the $N/O \propto O/H$ scaling dominate when the metallicity is above $\sim 1/3$ to $\sim 1/2$ solar (Shields 1976; Pagel Edmunds 1981; van Zee et 1998; Izotov Thuan 1999; Pettini et 2002). |
It has been noted (Παν et al. | It has been noted (Henry et al. |
2000: RKobuluicky Skillman 1996) that departures from the siuple ΑΟxOI velatiouship cau occur if the curichment is dominated by star formation in discrete bursts. | 2000; Kobulnicky Skillman 1996) that departures from the simple $N/O \propto O/H$ relationship can occur if the enrichment is dominated by star formation in discrete bursts. |
This situation leads to tinjc-dependent fluctuations in the N/O ratio because of differeut delays in the stellar release of N aud O ( ∙∙ | This situation leads to time-dependent fluctuations in the N/O ratio because of different delays in the stellar release of N and O (and C). |
the oven ]trend for increasing .3.5n with Ο/ΗHowever, remains. | However, the overall trend for increasing N/O with O/H remains. |
Moreover, here AIC of reports. fo our knowledge, of large N/O presented. ratios‘ in metalaA poor poor]iuterstellar "EUeas. | Moreover, there are no reports, to our knowledge, of large N/O ratios in metal poor interstellar gas. |
Laree N/O abundances are au iudicator of high metallicitiesCom cMiu ay scenario that involves a woellauixed interstellar mucin. | Large N/O abundances are an indicator of high metallicities in any scenario that involves a well-mixed interstellar medium. |
Early of investigationsthe abunudances in broad cuussion-line region (BELT) gas were based on several generally weak iuter-combination lines like NIVJALIS6, OT0AT663.. ΝΣΠΗΛΙΤΟΟ. aud i1A1909 (Shields 1976: Davidson 1977: Baldwin Netzer 1978: Osiner 1980: Gaskell et 11981: Uomoto 198D). | Early investigations of the abundances in broad emission-line region (BELR) gas were based on several generally weak inter-combination lines like $\lambda 1486$, $\lambda 1663$, $\lambda 1750$, and $\lambda 1909$ (Shields 1976; Davidson 1977; Baldwin Netzer 1978; Osmer 1980; Gaskell et 1981; Uomoto 1984). |
The results of these studies already indicated larger than solar iictallicity for the BELR eas. | The results of these studies already indicated larger than solar metallicity for the BELR gas. |
Receut studies of Ligh-redshift quasars (2 2.3) provide evidence for significantly eulianced uuctallicities wp to several times solar. | Recent studies of high-redshift quasars $z\ga 3$ ) provide evidence for significantly enhanced metallicities up to several times solar. |
These results obtained by studying the cussion liuc of propertiesquasars (Tamann Ferland 11993: Ferland et 11996: Dietrich ct 22002a: Dictrich Wilhelu-Erkeus 2000: Tamaun ct 22002: Warnerofthe et 22002) have been corroborated by studies intrinsic absorption lines (Petitjean et 11991: Moller et 11991: Taamaun 1997: Pettini 1999). | These results obtained by studying the emission line properties of quasars (Hamann Ferland 1993; Ferland et 1996; Dietrich et 2002a; Dietrich Wilhelm-Erkens 2000; Hamann et 2002; Warner et 2002) have been corroborated by studies of the intrinsic absorption lines (Petitjean et 1994; ller et 1994; Hamann 1997; Pettini 1999). |
The derived ligh chemical abunudanuces require an cra of major star formation at some earlier Uamann Ferland 11993) aud Ferlaucl et ((1996) show that omission Tine ratios involving vAI210 are particularly valuable. | The derived high chemical abundances require an era of major star formation at some earlier Hamann Ferland 1993) and Ferland et (1996) show that emission line ratios involving $\lambda 1240$ are particularly valuable. |
Generally. it is observed that vA1L240 is stronecr than expected iu the spectra of high redshift quasars compared to standard photoionization models with solar abundance. | Generally, it is observed that $\lambda 1240$ is stronger than expected in the spectra of high redshift quasars compared to standard photoionization models with solar abundance. |
Assuming uitroseu scales roughly as Z? at solar aud higher metallicities ILunaun Ferland 11993) sueeest vAI210/C A1519 and vAT210/IIe 11À1610 as valuable metallicity indicators. | Assuming nitrogen scales roughly as $Z^2$ at solar and higher metallicities Hamann Ferland 1993) suggest $\lambda 1240$ $\lambda 1549$ and $\lambda 1240$ $\lambda 1640$ as valuable metallicity indicators. |
Recently. Taman et (2002) preseuted results of a detailed investigation ou the iufluence of the photoiouizius continuum fux aud spectral shape. density. and metallicity on enüsson line ratios. | Recently, Hamann et (2002) presented results of a detailed investigation on the influence of the photoionizing continuum flux and spectral shape, density, and metallicity on emission line ratios. |
They further quautified the metallicity and ΑΠXZ? depeudence of various line ratios. including several weak inter-combinatiou lines. | They further quantified the metallicity and $N/H \propto Z^2$ dependence of various line ratios, including several weak inter-combination lines. |
They favor Naün]Al750/O WAL663 aud NVAI210/€O VIALO3L| C'IVAT519) line ratios as the most robust indicators to measure the gas chemical composition. | They favor $\lambda 1750$ $\lambda 1663$ and $\lambda 1240$ $\lambda 1034\, +$ $\lambda 1549$ ) line ratios as the most robust indicators to measure the gas chemical composition. |
Thsection 2section 2.we describewe thedeseribe thesample ofs;∖⊳∖⋅ the zz (ancl studied here. | In section 2 we describe the sample of the $z\ga 3.5$ quasars studied here. |
In section 3 the results N/Q the analvsis of the enüssion line spectra are no | In section 3 the results of the analysis of the emission line spectra are presented. |
We estimate the elemental abundance of. the line: enüttiug ⋅⋅↜↜eas based on several diagnostic∙↜⋅ enissionDoc line: ratios⋅⋅↴ (Ilzuuaun et ⊲⊔⊲≻22002). | We estimate the elemental abundance of the line emitting gas based on several diagnostic emission line ratios (Hamann et 2002). |
Usingn these ouission-line ratios. the mean metallicity is ZÍZ.cItoh for the BELR gas of the quasar sunple we observed. | Using these emission-line ratios, the mean metallicity is $Z/Z_\odot \simeq 4 \,{\rm to}\, 5$ for the BELR gas of the quasar sample we observed. |
The results are discussed aud compared with previous studies (e.9.. Ferland et 11996: IEuunaun Ferland 1999: Dietrich et 22002a: Dietrich. Wilhelu-Erkens 2000: Tlamann ct 22002: Warner et 22002) iu section [. | The results are discussed and compared with previous studies (e.g., Ferland et 1996; Hamann Ferland 1999; Dietrich et 2002a; Dietrich Wilhelm-Erkens 2000; Hamann et 2002; Warner et 2002) in section 4. |
The chemical composition of the DELR eas provides further evidence that the first episodes of major star formation started at a redshift of 54c GtoS. corresponding to an age of the universe of several LOS years. | The chemical composition of the BELR gas provides further evidence that the first episodes of major star formation started at a redshift of $z_f \simeq 6 \,{\rm to}\, 8$ , corresponding to an age of the universe of several $10^8$ years. |
This result is iu good agreement with recent model predictions relating quasay activity with the formation of massive spheroidal svstenis. e.g. the progenitors | This result is in good agreement with recent model predictions relating quasar activity with the formation of massive spheroidal systems, e.g., the progenitors |
and luminosity evolution models appearing at the faintest Hux levels (remembering that evolution is also. present in the other more established. components of the cppltldt mocel). | and luminosity evolution models appearing at the faintest flux levels (remembering that evolution is also present in the other more established components of the cppRR model). |
Serjeant et al. | Serjeant et al. |
(2000) found that the 151m integral counts could. indeed. be well fitted by a variety of evolving models (eppltlt. Franceschini et (1994)... Guideroni ct al. (1997). | \shortcite{serjeant00} found that the $\umu$ m integral counts could indeed be well fitted by a variety of evolving models (cppRR, Franceschini et \shortcite{fran94}, Guideroni et al. \shortcite{guid97}, |
. Xu et ab(1998))). | Xu et \shortcite{xu98}) ). |
Llowever. examining the 15jum differential counts provides à more cliscriminating picture. | However, examining the $\umu$ m differential counts provides a more discriminating picture. |
The dillerential counts clfeetively measure the slope of he integral counts and are hence more sensitive to subtle changes that may not be prominent in the integral counts at first sight. | The differential counts effectively measure the slope of the integral counts and are hence more sensitive to subtle changes that may not be prominent in the integral counts at first sight. |
Here. again we see a discrepancy between the wo evolution scenarios with the density evolution mocel oducing slightly higher dillerential. counts (i.e. steeper integral counts) at the faintest lluxes (perhaps hinting at »ing meareinally inconsistent with the observed. counts due to too much power at. higher redshifts?). | Here, again we see a discrepancy between the two evolution scenarios with the density evolution model producing slightly higher differential counts (i.e. steeper integral counts) at the faintest fluxes (perhaps hinting at being marginally inconsistent with the observed counts due to too much power at higher redshifts?). |
However. what is more striking is that neither the classical density nor luminosity evolution models can reproduce the subtle changes of slope in the Ίσμιαι counts that are not easily apparent [rom the integral source counts alone. | However, what is more striking is that neither the classical density nor luminosity evolution models can reproduce the subtle changes of slope in the $\umu$ m counts that are not easily apparent from the integral source counts alone. |
lt is apparent from the dilferential counts that the counts steepen drastically from ~3mJv-O4mdy (gldNids)=algsiam 3) above Euclidean expectations. turn over at —O0.4mJv and rapidly Datten at fainter Fluxes. | It is apparent from the differential counts that the counts steepen drastically from $\sim$ 3mJy-0.4mJy $lg(dN/dS)= \alpha lgS, \alpha \approx -3$ ) above Euclidean expectations, turn over at $\sim$ 0.4mJy and rapidly flatten at fainter fluxes. |
The analysis by Serjeant et al. | The analysis by Serjeant et al. |
(2000) concluded: that models. incorporating pure luminosity evolution provided a better fit than the density evolving mocdels although they assumed a lower value for the density evolution anc moreover their observations cid not extend. down in detail below 10m.Jy. in the dillerential Counts. | \shortcite{serjeant00} concluded that models incorporating pure luminosity evolution provided a better fit than the density evolving models although they assumed a lower value for the density evolution and moreover their observations did not extend down in detail below 10mJy in the differential counts. |
At 601un the LAS data does not really extend to deep enough Illuxes to gain meaningful insight [rom the mocels. with both the luminosity and density evolution. scenarios fitting both the integral and cilferential counts. although the faintest dilferential counts would imply stronger evolution than currently assumed. | At $\umu$ m the IRAS data does not really extend to deep enough fluxes to gain meaningful insight from the models, with both the luminosity and density evolution scenarios fitting both the integral and differential counts, although the faintest differential counts would imply stronger evolution than currently assumed. |
However. at 17015 none of the models can account for the observed. source counts at fainter κος (although the brightest points are fitted by all models). | However, at $\umu$ m none of the models can account for the observed source counts at fainter fluxes (although the brightest points are fitted by all models). |
At 170pum we could envision a general increase in the magnitude of the evolution as generally [itting the observed. counts (Le. increasing the density evolution power law index 9). | At $\umu$ m we could envision a general increase in the magnitude of the evolution as generally fitting the observed counts (i.e. increasing the density evolution power law index $g$ ). |
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