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We saw (that retrograde precession occurs. if. the secondary star in v-Octantis is in turn a binary system inclined more than 45° with respect to the main hinary’s orbit (Fig.	We saw that retrograde precession occurs if the secondary star in $\nu$ -Octantis is in turn a binary system inclined more than $45^\circ$ with respect to the main binary's orbit (Fig.
2).	2).
In this scenario. we saw that a periodogram of the residuals. leftover from. fitting a Weplerian orbit to the main binary could exhibit peaks that might be mistaken as planets.	In this scenario, we saw that a periodogram of the residuals leftover from fitting a Keplerian orbit to the main binary could exhibit peaks that might be mistaken as planets.
We saw that these peaks appeared close to harmonies of the main binary’s orbital frequenev.	We saw that these peaks appeared close to harmonics of the main binary's orbital frequency.
In particular. we showed an example (Sect.	In particular, we showed an example (Sect.
4.3 and Fie.	4.3 and Fig.
4) where the fake planet's period is nearly commoensurate (ratio 2/5) with the main binary’s period. which is exactly. what happens in z-Octantis.	4) where the fake planet's period is nearly commensurate (ratio 2/5) with the main binary's period which is exactly what happens in $\nu$ -Octantis.
Therefore. we propose that a hidden binary system could. mimic a planct similar to the one reported in #-Octantis by 2..	Therefore, we propose that a hidden binary system could mimic a planet similar to the one reported in $\nu$ -Octantis by \citet{Ramm_etal2009}.
In order to estimate inner binary parameters that lead to retrograde precession ofthe outer binary at arate c»= Ü.SG/gr. we set aj=a,0.35 AU (Iq. 34)).	In order to estimate inner binary parameters that lead to retrograde precession ofthe outer binary at arate $\dot{\omega}_2=-0.86^\circ/yr$ , we set $a_1=a_c=0.35$ AU (Eq. \ref{acritical}) ),
7=GO” and c,=0.76 tthe inner binary is at the Wozai stationary solution).	$i=60^\circ$ and $e_1=0.76$ the inner binary is at the Kozai stationary solution).
Replacing these in Eqs. (	Replacing these in Eqs. (
25) and (26) with a2= AU. mo=L4Al. and m,=0.5M.. we obtain estimates for the inner binarvs masses of my,=0.23M. and my=0.27M...	25) and (26) with $a_2=2.55$ AU, $m_2=1.4\,M_{\odot}$ and $m_b=0.5\,M_{\odot}$, we obtain estimates for the inner binary's masses of $m_1=0.23\,M_{\odot}$ and $m_0=0.27\,M_{\odot}$.
In Table 5 we show the results of fitting a precessing Weplerian orbit. (fit. 1) and a fixed. Ixeplerian orbit. (fit 0) to such triple system. (simulation D).	In Table 5 we show the results of fitting a precessing Keplerian orbit (fit 1) and a fixed Keplerian orbit (fit 0) to such triple system (simulation I).
We see that fit (1) is better than fit (0) and is comparable with the fit of à planet at 495 dav.	We see that fit (1) is better than fit (0) and is comparable with the fit of a planet at 495 day.
We can also estimate inner binary parameters by performing N-bocly fits to 7-Octantis radial velocity data.	We can also estimate inner binary parameters by performing N-body fits to $\nu$ -Octantis radial velocity data.
We assumed m»=L4Al.. and fixed the outer binary’s orbit at fo=10.85 and O»=ST which are the parameters inferred from. the spectroscopic-astrometric solution (Table 4).	We assumed $m_2=1.4\,M_{\odot}$, and fixed the outer binary's orbit at $I_2=70.8^\circ$ and $\Omega_2=87^\circ$ which are the parameters inferred from the spectroscopic-astrometric solution (Table 4).
We present the best fit solution in Table 6.	We present the best fit solution in Table 6.
Retrograde oecession occurs because 7=122°.	Retrograde precession occurs because $i=122^\circ$.
Moreover. the best πουν N-bocky fit Clable 6) is comparable to the Ixeplerian Xanet fit (Lable 4).	Moreover, the best binary N-body fit (Table 6) is comparable to the Keplerian planet fit (Table 4).
Although this solution is unstable oeause αι8]Loa. we can obtain “equivalent” stable configurations by reducing αι. while maintaining my,538M. constant. and increasing the ratio m/m». so that he Ixeplerian term and the quadrupole interaction term are th kept constant.	Although this solution is unstable because $a_1\approx 1.5\, a_c$, we can obtain "equivalent" stable configurations by reducing $a_1$, while maintaining $m_b=0.538\,M_{\odot}$ constant, and increasing the ratio $m_1/m_b$, so that the Keplerian term and the quadrupole interaction term are both kept constant.
We performed a simulation of such a stable configuration with a,=a,=0.35 AU. m,=0.109AL. and my=0.429M...	We performed a simulation of such a stable configuration with $a_1=a_c=0.35$ AU, $m_1=0.109\,M_{\odot}$ and $m_0=0.429\,M_{\odot}$.
In Table 5 we show the results of fitting a precessing Ixeplerian orbit (fit 1) and a fixed Ixeplerian orbit (it 0) to such triple system. (simulation LH).	In Table 5 we show the results of fitting a precessing Keplerian orbit (fit 1) and a fixed Keplerian orbit (fit 0) to such triple system (simulation II).
We see that fit (1) is better than fit (0) but fit (1) is better than the fit of a planet at 452 day.	We see that fit (1) is better than fit (0) but fit (1) is better than the fit of a planet at 452 day.
The values of the precession rate after lit (1). and both V and residuals (ims) after fitting the planets orbit are almost equal to those obtained for the real system /z-Octantis (ef.	The values of the precession rate after fit (1), and both $\sqrt{\chi^2}$ and residuals $rms$ ) after fitting the planet's orbit are almost equal to those obtained for the real system $\nu$ -Octantis (cf.
Table 4).	Table 4).
From Γαρίο 5. as described. above. we see that. [or simulated. data the Π of a precessing Weplerian orbit. is comparable (simulation E) or better (simulation IH) than the fit of a planet. while in the real case (v-Octantis) the fit of a planet is currently better than the fit of a precessing Keplerian orbit.	From Table 5, as described above, we see that for simulated data the fit of a precessing Keplerian orbit is comparable (simulation I) or better (simulation II) than the fit of a planet, while in the real case $\nu$ -Octantis) the fit of a planet is currently better than the fit of a precessing Keplerian orbit.
However. we stress that there are. many more combinations of parameters (mo. mj. 7 and ej) that	However, we stress that there are many more combinations of parameters $m_0$ , $m_1$ , $i$ and $e_1$ ) that
(3.6Alpe.Freedmanetal.1994).. (Peimbert&Torres-PeimbertSpevbroeck1982). (LowIonizationNuclearEmissionlineRegion:1996) fa (Peimbert&Torres-Peimbert1951).. (Lyc10M (Elvis&van1982:Fabbiano1988) E1.35 (Ishisakieta	\citep[3.6 Mpc,][]{fre}, \citep{pei,elvis} \citep[Low Ionization Nuclear Emission line Region;][]{ho} $H\alpha$ \citep{pei}, \citep{biet} $L_X \sim 10^{40}$ \citep{elvis,fab88} $\Gamma\sim1.85$ \citep{ishi,pelle}.
l.1996:Pellegrini2000).. (Ilo.Filippenko.Sargent1996) (Boweretal.2000) 3x10M.<Mpy«610*M.. Ishisakietal.1996)) (DL1.85). KT~0.86(Ishisakietal.1996)... Inml	\citep{ho} \citep{bower} $3\times 10^6M_{\odot}< M_{BH}<6\times 10^7M_{\odot}$ \citealp{ishi}) \citealp{pelle}) $\Gamma \sim 1.85$ $kT\sim 0.86$\citep{ishi}. \citet{imm},
er&Wang(2001).. T=1.85 Barretal(1985) (~2	$\Gamma=1.85$ $\sim600$ \citet{barr} $\sim2$
]t is generally accepted that galactic cosmic ravs are accelerated in supernova remnants.	It is generally accepted that galactic cosmic rays are accelerated in supernova remnants.
The process of dilfusive shock acceleration (Ixrvmskii1977:Axfordetal.Bell1978:Blandford&OstrikerLOTS) remains the most likely mechanism. for⋅ producing⋠‘ and maintainingMNu the observed spectrum.	The process of diffusive shock acceleration \citep{krymskii77,axfordetal77,bell78a,blandfordostriker78} remains the most likely mechanism for producing and maintaining the observed spectrum.
asThere is: now a growing. wealth of observational. evidence. supporting: this. scenario.	There is now a growing wealth of observational evidence supporting this scenario.
DoThe detection: of EETe gamma-ray emission. [rom nearby remnants. Confirms the presence of electrons. ancl possibly protons. with energies of llf ⋜⊔↓⋖⋅⋜↧⊳∖↿↓∪⋖⊾∖↿∖⋖⋅⊳⋏∙≟⊳∐↓⊔⋯⊔⊾∖↽∐∪⇂⊔↓⋜⋯⊔−≽∪∪≤∩↣↓	The detection of TeV gamma-ray emission from nearby remnants confirms the presence of electrons, and possibly protons, with energies of at least $10^{14}$ eV \citep[e.g.][]{hintonhofmann09}.
⊔⋯∐∐↓∩⊔⋡:⋅− ↔ high. resolution. observations. of⋅ narrow non-thermal X-ray. filaments⋅ at the outer shocks of⋅ several shell-type supernova remnants favour⋅ a model in. which. the highest. energy electrons are produced. directly.. at the shock. consistent. with. the predictionsD. of ⋅⊀⋅∢dilfusive shock. acceleration.⊀ (e.g.nne 2007).	In addition, high resolution observations of narrow non-thermal X-ray filaments at the outer shocks of several shell-type supernova remnants favour a model in which the highest energy electrons are produced directly at the shock, consistent with the predictions of diffusive shock acceleration \citep[e.g.][]{vinklaming03,bambaetal05,uchiyamaetal07}.
. cpThese filaments⋅ also provide. evidence: for⋅ strong ⊀⋅ ∢⋅⊀ . ↔⊀ ⋅ ⊔↓⋜↧⋏∙≟⊔∢⊾⇂⊔⇍∐∢⊾↓∠⇂⋜⋯↓↓≻↓↓↓∐∼⋜∐↓∪⊔↓⊔⇂↓↥⋖⊾∖⇁	These filaments also provide evidence for strong magnetic field amplification in the vicinity of the shock.
⊔⇍↓⊔↓↿∙∖⇁∪⇂↿↓⊔⋅⊳∖↓↥⋯∙↓∡⋡ .generation of⋅ strong. magnetic turbulenceMN is vital ⋅for ⋅& ⋜⋯∠⇂⋜↧∣⋡∪∖⇁∢⊾↿∖∟⋜↧⋏∙≟⋜↧⋏∙≟∢⋅⊾∖↽≺⊲∢⊾⊳∖⋜⊔⋅⊳∖↳∖⇁↓≤⋗↖∖∶∫≻∶≧⋖⊾∐⊾∖↽∟⋯⇍⋖⊾↓∡⇉∪∪↓⊐↦	The generation of strong magnetic turbulence is vital for the acceleration of cosmic rays to the knee $\sim10^{15.5}$ eV) and above \citep{lagagecesarsky83, belllucek01}.
While several mechanisms for amplifving magnetic fieles to values in excess of the shock compressed interstellar fields⋅ have been proposed. those that result in. the transfer⋅ of . . ⊔↓≻⊳∖↿↓⋅⋖⋅⋜⋯↓≼⇍∪⊳∖⊔↓⊔⇍−↓⋅⋜↧∙∖⇁⊳∖⇂↓⋅⋖⋅⋜⋯⊔⊔⋏∙≟⋖⋅⊔⋖⋅↓⋅⋏∙≟∙∖⇁⋯↿↥⋖⋅⊔↓⋜↧⋏∙≟⊔∢⊾⇂⊔⇍ ∢⋅field are of⋅ greatest relevance ⋅for diffusiveMM shock acceleration.	While several mechanisms for amplifying magnetic fields to values in excess of the shock compressed interstellar fields have been proposed, those that result in the transfer of upstream cosmic-ray streaming energy to the magnetic field are of greatest relevance for diffusive shock acceleration.
. opsThe non-resonant ποσο first⋅ identified.» by Bell(2004) has been demonstrated: to grow rapidly.. however. the characteristic2. wavelength of the amplifiecl field is. predominantly. on a length. scale shorter than the &vroradius of the driving particles.	The non-resonant mode first identified by \cite{bell04} has been demonstrated to grow rapidly, however, the characteristic wavelength of the amplified field is predominantly on a length scale shorter than the gyroradius of the driving particles.
Under certain conditions. other short-wavelength instabilities may dominate (c.gBret2009:Riquelme& 2011)..	Under certain conditions, other short-wavelength instabilities may dominate \citep[e.g][]{bret09,riquelmespitkovsky10,lemoinepelletier10,nakaretal11}.
While such instabilities may be sullicient to explain the [large magnetic field. values inferred. [rom observations. in. order to facilitate⋅↔ rapid. acceleration. to higher. energies... it is. necessaryu to generate. field⋅ structure∙ on scales comparable∙ with. the gvroradius. of⋅ the highest. energy. particles. (Bell>&Lucek2001:Revilleetal. 2008).	While such instabilities may be sufficient to explain the large magnetic field values inferred from observations, in order to facilitate rapid acceleration to higher energies, it is necessary to generate field structure on scales comparable with the gyroradius of the highest energy particles \citep{belllucek01,revilleetal08}.
. The generation of large scale field structures has been investigated in the context of filaments. or beams. in Dellomar ↿∖−≽∪∪⋅≻∃↦∖∖⋎↓↕⋖⊾↓⋅∢⋅⋜↧↓≻↓⋅⋖⋅⊣⊾⇀∖↓⊳∖⊔⊔⋏∙≟↓≻↓⋅∪∐↓⋖⊾⇂∪↓⋅⇂↓↕∢⊾≼∼∪⊳∖⊔↓⊔∼↓⋅⋜↧∙∖⇁Do. ⋅⋅ ⊀ distribution⋠⋠⋠ was assumed.	The generation of large scale field structures has been investigated in the context of filaments, or beams, in \citet{bell05}, where a pre-existing profile for the cosmic ray distribution was assumed.
In this. paper. we demonstrate that the distribution⊲⊀⊀ of⋅ relativistic.. particles: is 2.inherently κ. ⊔∪⊔−⊔⊔↓⇂∪↓⋅⊔↓⊳⋜⋯∠⇂⇂↓⋯↿∐↓⋜⋯↓⋖⋅⊔↿⋜⊔↓∪⊔∪⊓⇍⊔↓⋅⊳∖⋜↧⊳∖⋜↧⊔⋜∐⊔↓⋅⋜↧↓⋅. consequence» off cosmic-ray. streaming.	In this paper, we demonstrate that the distribution of relativistic particles is inherently non-uniform, and that filamentation occurs as a natural consequence of cosmic-ray streaming.
.” A\ similar »phenomenonpl occurs in laser plasmas whereby photon beams filament due ↿∪↿↓↥∢⊾↓⋅⊔⋯⇂⊳∖⋖⋅∐⋅−⇂⋅⋯∼⊔⊳∖⊲↓⊔⋃↕↓↕∢⋅⇀∖≻⋜↧↓↥∠⇂↕↓↕∙≱≼⇍⋜↧∖⇁↕⇂↕⋖⋅≻↿∖∢⊾⊳⋃⊳≺⊲↓⋅⋜∟∖⋯⊔UU .MeCrorv1984).	A similar phenomenon occurs in laser plasmas whereby photon beams filament due to thermal self-focusing in expanding cavities \cite[e.g.][]{craxtonmccrory84}.
.27 We show .here that in the case of cosmic- Phe.. this process results in the growth of magnetic field on laree scales.	We show here that in the case of cosmic-rays, this process results in the growth of magnetic field on large scales.
The development of the filamentation and larec-scale field is investigated. analytically in à two dimensional slab symmetric geometry. and verified using hybrid. particle-AID simulations.	The development of the filamentation and large-scale field is investigated analytically in a two dimensional slab symmetric geometry, and verified using hybrid particle-MHD simulations.
. mThe non-linear. feedback.⋅ between: ultra Dl . ↓⋅∢⊾↓⋜∐↓∖⇁↓⊳∖∣⊔∼↓≻⋜⊔⋅↿⊔⇍↓⋖⋅⊳∖⋜↧⊔∠⇂↿↓↕⋖⋅∣⋡⋯∙↳⋏∙≟↓⋅∪⊔⊔∠⇂↓≻↓⋜↧⊳∖⊔⋯⊳⋜⋯∠↿↓⊔⋅ resulting DOlarge-scaleT fields will have\| importantpoms implicationspuc for the acceleration of cosmic ravs to energies? above the knee in. supernova remnants. and also their. escape.	The non-linear feedback between ultra relativistic particles and the background plasma, and the resulting large-scale fields will have important implications for the acceleration of cosmic rays to energies above the knee in supernova remnants, and also their escape.
The outline of the paper is as follows.	The outline of the paper is as follows.
In the next section we develop the analytic model that describes the cosmic-ray filamentation.	In the next section we develop the analytic model that describes the cosmic-ray filamentation.
It ids clemonstrated that this introduces a long wavelength: instability in the precursors	It is demonstrated that this introduces a long wavelength instability in the precursors
oorbit [ound by lindicated that the finding is not highly significant.	orbit found by indicated that the finding is not highly significant.
Pourbaix (2001). who studied the significance of the derived [Hipparcos orbits of 42 stars that host planet candidates. [ound that the statistical significance of the oorbit is al the level.	Pourbaix (2001), who studied the significance of the derived Hipparcos orbits of 42 stars that host planet candidates, found that the statistical significance of the orbit is at the level.
He used. an F-test to evaluate the improvement of the fit to the IHlipparcos data resulting from the additional parameters of the binary orbit.	He used an F-test to evaluate the improvement of the fit to the Hipparcos data resulting from the additional parameters of the binary orbit.
Hle also conclIunded that the other orbits of anre statisticallv non significant.	He also concluded that the other orbits of are statistically non significant.
The use of the F-distribution assumes Gaussianity of the individual measuremenis. an assumption (hat might not be well justified for the Iipparcos data.	The use of the F-distribution assumes Gaussianity of the individual measurements, an assumption that might not be well justified for the Hipparcos data.
Arenou et ((1995) proved that the Hipparceos parallaxes and zero-points are Gaussian. but thev do not analvze the distribution of the individual measurements of each star.	Arenou et (1995) proved that the Hipparcos parallaxes and zero-points are Gaussian, but they do not analyze the distribution of the individual measurements of each star.
We (Zucker Mazel 2001) avoided the assumption of Gaussianitv by using (he permutation test. which belongs to the class of distribution-Iree tests (e.g. Good 1994) and thus is more robust against modeling problems of the measurement process.	We (Zucker Mazeh 2001) avoided the assumption of Gaussianity by using the permutation test, which belongs to the class of distribution-free tests (e.g. Good 1994) and thus is more robust against modeling problems of the measurement process.
We performed (he permutation test by generating simulated data [rom (hie verv same astrometric measurements ofCrD.	We performed the permutation test by generating simulated data from the very same astrometric measurements of.
. If there was some evidence of an orbit in the measurements. it should be ruined by the permutation. aud thus no random permutations would vield a comparable orbit.	If there was some evidence of an orbit in the measurements, it should be ruined by the permutation, and thus no random permutations would yield a comparable orbit.
However. if the derived orbit was only spurious. some random permutations should be able to reproduce a similar elfect.	However, if the derived orbit was only spurious, some random permutations should be able to reproduce a similar effect.
In a sense. we let the data “speak for themselves" and do not have to assume any specific distribution for the measurements or the errors.	In a sense, we let the data “speak for themselves” and do not have to assume any specific distribution for the measurements or the errors.
We found that the significance of the astrometric solution of iis ad the level.	We found that the significance of the astrometric solution of is at the level.
The study of the significance of the aastrometric detection led us to check carefullyaff (he announced extrasolar planet ancl brown dwarf candidates with available data.	The study of the significance of the astrometric detection led us to check carefully the announced extrasolar planet and brown dwarf candidates with available data.
All together we present here an analysis of 47 planet candidates and. 14 brown cdwarls.	All together we present here an analysis of 47 planet candidates and 14 brown dwarfs.
We first derive from the Ilipparcos data the best astromeltric orbit and (hen test unequivocally ils significance. wilh an approach which is [ree of any assumptions about the errors of the Lipparecos measurements.	We first derive from the Hipparcos data the best astrometric orbit and then test unequivocally its significance, with an approach which is free of any assumptions about the errors of the Hipparcos measurements.
All the derived orbils of the stars that harbor extrasolar planets turned out to be insignificant.	All the derived orbits of the stars that harbor extrasolar planets turned out to be insignificant.
Pourbaix and Arenou (2001) reached similar conclusions.	Pourbaix and Arenou (2001) reached similar conclusions.
Although the derived orbits are insignilicant. Eve nevertheless used them to derive upper limits for the corresponding astrometric motions and for (he masses of the unseen companions.	Although the derived orbits are insignificant, we nevertheless used them to derive upper limits for the corresponding astrometric motions and for the masses of the unseen companions.
The previous study. which derived upper limits for planet candidates (Perrvinan οἱ 11996) searched the IHipparcos data for any. possible astrometric orbital periodicity [or each of the three stars considered then.	The previous study which derived upper limits for planet candidates (Perryman et 1996) searched the Hipparcos data for any possible astrometric orbital periodicity for each of the three stars considered then.
They derived a periodgram Irom (he corresponding astromeltric aanplitudes aud obtained an upper limit for the stellar astrometrie motion from	They derived a periodgram from the corresponding astrometric amplitudes and obtained an upper limit for the stellar astrometric motion from
large scale fluctuaions with amplitude of <12 over the whole profile. with ouly two exceptions reaching locally values ofC.	large scale fluctuations with amplitude of $\leq 1 - 2$ over the whole profile, with only two exceptions reaching locally values of.
This implies that overall a Nuker law reproduces quiο well the ooOalaxies profiles.	This implies that overall a Nuker law reproduces quite well the galaxies profiles.
The parameters derived from the Nuker fits are reported in Tab.	The parameters derived from the Nuker fits are reported in Tab.
3 and L.	\ref{tabobs1} and \ref{tabobs2}.
The brightuess at the break radius ji, has beei converted. to the standard C'ousiu-Jolson system aid corrected for Calactic extinction.	The brightness at the break radius $\mu_b$ has been converted to the standard Cousin-Johnson system and corrected for Galactic extinction.
Break. radii are in arcsecouds.	Break radii are in arcseconds.
Sources are separated iuto power-law aud core galaxies adopting the stareard values. nes 5x0.3 are classufed as core-galaxies. while 5=0.5 are power-law galaxies.	Sources are separated into power-law and core galaxies adopting the standard values, i.e. $\gamma \leq0.3$ are classified as core-galaxies, while $\gamma \geq0.5$ are power-law galaxies.
We also considere as power-law galaxies. following the scheme of Lauereal.(1995).. all those objects (6) iu which uo break in the brigltuess profile is seen at the UST resolution limit.	We also considered as power-law galaxies, following the scheme of \citet{lauer95}, all those objects (6) in which no break in the brightness profile is seen at the HST resolution limit.
More specifically. we included ii this class the ealaxies with rjx072, as the region below the break is not sufficiently sampled (Gehen uot simply muresolyed) to provide au accurate estimate of 5.	More specifically, we included in this class the galaxies with $r_b \leq 0\farcs2$, as the region below the break is not sufficiently sampled (when not simply unresolved) to provide an accurate estimate of $\gamma$.
Poteutially. these objects might be core galaxies located at a sufficiently large distance so that a compact shallow core caunot be resolved.	Potentially, these objects might be core galaxies located at a sufficiently large distance so that a compact shallow core cannot be resolved.
Iudeed. there is oue core-galaxy iu our sauple. UCC 7760 (NGC. 1552, NL 89) with v=392 kins J\|. with a well resolved core. kj=0719.y corresponding to 13 pe. which would have been nüsclassified with this scheme if it was located at a larger distance.	Indeed, there is one core-galaxy in our sample, UGC 7760 (NGC 4552, M 89) with v=392 km $^{-1}$, with a well resolved core, $r_b = 0\farcs49$ corresponding to 13 pc, which would have been misclassified with this scheme if it was located at a larger distance.
Thus we prefer in the roinaiuiug of the paper to mark this class differeutlv from objects with well established power-law profiles.	Thus we prefer in the remaining of the paper to mark this class differently from objects with well established power-law profiles.
Iu the	In the
spectrin of the Bl star forming knot iu NCCI1711 (Fie.	spectrum of the B1 star forming knot in NGC1741 (Fig.
5 top) (Conti. Leitherer Vacca 1996)).	\ref{Fig5} top) (Conti, Leitherer Vacca \cite{conti96}) ).
We rejected those features such that: a) the spectral profile was (taking errors into account) narrower than the instrumental profile (IP) (2.98 in the rest frame) and/or b) the absorbed flux was lower than the detection limit.	We rejected those features such that: a) the spectral profile was (taking errors into account) narrower than the instrumental profile (IP) (2.98 in the rest frame) and/or b) the absorbed flux was lower than the detection limit.
This was the case of SIITAT260. L. OVAISTI. SillIALIT7. SVALSO2.	This was the case of $\lambda$ 1260.4, $\lambda$ 1371, $\lambda$ 1417, $\lambda$ 1502.
Tn order to calculate the detection limit for an absorption feature at a eiveu position. we created Catssians with the expected EWIIM (IP in all cases) anc varied the amplitude (the profiles could be broader. but this jus nieans that a larger flux would be needed for detection).	In order to calculate the detection limit for an absorption feature at a given position, we created Gaussians with the expected FWHM (IP in all cases) and varied the amplitude (the profiles could be broader, but this just means that a larger flux would be needed for detection).
The Gaussians were added to the coutinuun near the expected position.	The Gaussians were added to the continuum near the expected position.
The upper limit was chosen by eve. as the flux of that Cassiani that we considered detectable.	The upper limit was chosen by eye, as the flux of that Gaussian that we considered detectable.
We present in Fie.	We present in Fig.
6 the fits to those features that we accepted as real.	\ref{Fig6} the fits to those features that we accepted as real.
Except for CIV (for which the original spectrum is shown). we prescut sinoothed spectra to make the figures clearer.	Except for CIV (for which the original spectrum is shown), we present smoothed spectra to make the figures clearer.
The fits were done to the original (non-simoothed) spectra.	The fits were done to the original (non-smoothed) spectra.
There are some sky cussion residuals ou the blue side of the CTV absorption feature. but they do not affect the fit (see Fie. D).	There are some sky emission residuals on the blue side of the CIV absorption feature, but they do not affect the fit (see Fig. \ref{Fig4}) ).
We present in Table 2 sone paraieters obtained frou the fits: wavelength. line ideutification. EW (rest frame) and FWIOA	We present in Table 2 some parameters obtained from the fits: wavelength, line identification, EW (rest frame) and FWHM.
L The values measured in the racio galaxy LC11.17 (2 =3.80) (Dev et al.	The values measured in the radio galaxy 4C41.17 $z=$ 3.80) (Dey et al.
1997) aud the associated absorption svstei ofthe quasar 3CI9L (2 =1.95) (Balicall. Sargeut Schinidt 1967 1967)) are also shown.	1997) and the associated absorption system of the quasar 3C191 $z=$ 1.95) (Bahcall, Sargent Schmidt 1967 \cite{bahc67}) ) are also shown.
The fitting to SUVAALS93.1102 is dificult due to the low S/N ratio. however. the presence of two P-Cyenui profiles for the A1393.8 and SiTVAL102.8 lines ls sugeestte w the data.	The fitting to $\lambda\lambda$ 1393,1402 is difficult due to the low S/N ratio, however, the presence of two P-Cygni profiles for the $\lambda$ 1393.8 and $\lambda$ 1402.8 lines is ted by the data.
The best fit is obtained bv considering two absorption aud two emissiou features;	The best fit is obtained by considering two absorption and two emission features.

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